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1 2 Text and Left image from the CISM Summer School (Boulder, August 2013) SW101_4_Flares https://www.bu.edu/cism/SummerSchool/summerlist.html Right animation from ESA: http://sci.esa.int/cluster/36447-direct-observation-of-3d-magnetic-

  1. Shibata et al. (1995): Hot-Plasma Ejections Associated with Compact-Loop Solar Flares http://adsabs.harvard.edu/abs/1995ApJ...451L..83S One of the biggest discoveries by the soft X-ray telescope (SXT) (Tsuneta et al. 1991) aboard Yohkoh (Ogawara et a1. 1991) is that of cusp-shaped loop structures in long duration event (LDE) flares (Tsuneta et al. 1992a) and large-scale arcade loops associated with filament eruption or coronal mass ejections (Tsuneta et al. 1992b; Hanaoka et al. 1994; McAllister et at. 1995; Hudson, Haisch, & Strong 1995). The observed loop configurations of the LDE flares and arcade loops are quite similar to the magnetic field configuration suggested by the classical two-ribbon flare model (Carmichael 1964; Sturrock 1966; Hirayama 1974; Kopp & Pneuman 1976). This model, which is hereafter called the CSHKP model, predicts that magnetic fields are first opened up by global MHD instability associated with filament eruption to form vertical current sheet, and then magnetic field lines in the current sheet successively reconnect to form apparently growing flare loops. Recently, using the hard X-ray telescope (HXT) (Kosugi et al. 1991) aboard Yohkoh, Masuda et al. (1994) discovered that some of impulsive compact-loop flares occurring near the solar limb, a loop top hard X-ray (HXR) source appeared well above a soft X-ray (SXR) bright loop during the impulsive phase. This indicates that the impulsive energy release did not occur within the soft X-ray loop but above the loop. This is a quite exciting discovery because bright soft X-ray loops were often considered to be evidence of “loop flares” in which energy release occurs within the loop, as discussed above. One possible physical mechanism to produce such loop top hard X-ray source is magnetic reconnection occurring above the loop; i.e., a high-speed jet is created through the reconnection and collides with the loop top, producing fast-mode MHD shock, superhot plasma, and/or high- energy electrons emitting hard X-rays. In this sense, the discovery of the loop-top HXR source may open a possibility to unify two distinct classes of flares, two-ribbon flares (or LDE flare) and compact-loop flares (or impulsive flare), by the single mechanism of magnetic reconnection (Shibata 1995). Also: http://solar.physics.montana.edu/magara/Research/Topics/cshkp.html 11

  2. Martens and Kuin (1989): A circuit model for filament eruptions and two-ribbon flares http://adsabs.harvard.edu/abs/1989SoPh..122..263M Shibata et al. (2011): Solar Flares: Magnetohydrodynamic Processes http://solarphysics.livingreviews.org/Articles/lrsp-2011-6/ Several classic models based on magnetic reconnection have been proposed to explain the phenomenological aspect of flares: Carmichael (1964), Sturrock (1966), Hirayama (1974), and Kopp and Pneuman (1976) (see Figure 2). These models assume more or less a similar configuration of magnetic field and its dynamic process, so these models are called with a single name, CSHKP model (ˇ Svestka and Cliver, 1992; Sturrock, 1992; Shibata, 1999). The CSHKP model has been a standard model of flares, and the basic features of this model are explained in Figure 3. Shibata (1999): Evidence of Magnetic Reconnection in Solar Flares and a Unified Model of Flares http://adsabs.harvard.edu/abs/1999Ap%26SS.264..129S Shibata (1994) showed that, because there was an HXR on top of the SXR loop tops. As the HXR are produced by high energy electrons, this means that the main energy release took place outside/above the SXR loops. Also: http://solar.physics.montana.edu/magara/Research/Topics/cshkp.html 12

  3. Figure source: Benz, A. O., 2008 Living Reviews: http://solarphysics.livingreviews.org/Articles/lrsp-2008-1/ 1.3 The phases of flares The timing of the different emissions of the same flare is presented schematically in Figure 2. In the pre-flare phase the coronal plasma in the flare region slowly heats up and is visible in soft X-rays and EUV. A large number of energetic electrons (up to 10 38 ) and ions (with similar total energy) is accelerated in the impulsive phase, when most of the energy is released. The appearance of hard X-ray footpoint sources at chromospheric altitude is a characteristic of this phase (Hoyng et al.,1981). Some high-energy particles are trapped and produce intensive emissions in the radio band. The thermal soft X- ray and Hα emissions finally reach their maxima after the impulsive phase, when energy is more gently released, manifest in decimetric pulsations, and further distributed. The rapid increase in Hα intensity and line width has been termed flash phase. It coincides largely with the impulsive phase, although Hα may peak later. In the decay phase, the coronal plasma returns nearly to its original state, except in the high corona (>1.2R, where R is the photospheric radius), where magnetic reconfiguration, plasma ejections and shock waves continue to accelerate particles, causing meter wave radio bursts and interplanetary particle events. [Kane, 1974: The gradual phase consists of the flash phase and the decay phase. Here, the flash phase concerns the energy release as the sudden brightening. (Fundamentals of Solar astronomy, Figure 5.40).] 13

  4. Top: From Hanaoka et al., 1999: Radio and X-ray Observations of the Flares Caused by Interacting Loops http://solar.nro.nao.ac.jp/meeting/nbym98/PDF/hanaoka_2.pdf And Bottom: From Maria Massi (What is a solar flare?) http://www3.mpifr-bonn.mpg.de/staff/mmassi/#coronae1 Other example: Solar flare mechanism: http://www.stce.be/news/265/welcome.html 14

  5. Magnetic flux emergence: X6.9 flare on 9 August 2011: http://www.stce.be/news/353/welcome.html Blue/black is negative (inward) magnetic polarity, red/white is positive (outward) polarity Helical energy storage: X2.2 flare on 15 February 2011 Velareddi et al. (2012): On the role of rotating sunspots in the activity of solar active region NOAA 11158 http://iopscience.iop.org/article/10.1088/0004-637X/761/1/60/pdf Jiang et al. (2011): Rapid sunspot rotation associated with the X2.2 flare on 2011 February 15 http://iopscience.iop.org/article/10.1088/0004-637X/744/1/50 Also at PhysOrg: https://phys.org/news/2011-04-rotating-sunspots-super-solar-flare.html 15

  6. Magnetic flux emergence: X6.9 flare on 9 August 2011: http://www.stce.be/news/353/welcome.html Helical energy storage: X2.2 flare on 15 February 2011 Velareddi et al. (2012): On the role of rotating sunspots in the activity of solar active region NOAA 11158 http://iopscience.iop.org/article/10.1088/0004-637X/761/1/60/pdf Jiang et al. (2011): Rapid sunspot rotation associated with the X2.2 flare on 2011 February 15 http://iopscience.iop.org/article/10.1088/0004-637X/744/1/50 Also at PhysOrg: https://phys.org/news/2011-04-rotating-sunspots-super-solar-flare.html 16

  7. Kink instability Török et al. (2010): The writhe of helical structures in the solar corona http://www.aanda.org/articles/aa/full_html/2010/08/aa13578-09/aa13578-09.html http://www.lmsal.com/TRACE/POD/TRACEpodarchive14.html#movie61 (27 May 2002; M2 ; NOAA 9957) Unstable if twist ~2.5pi (Török et al., 2003: http://www.aanda.org/articles/aa/pdf/2003/30/aah4206.pdf ). Unstable magnetic fields Collateral damage: http://www.stce.be/news/361/welcome.html Shen et al. (2014): A Chain of Winking (Oscillating) Filaments Triggered by an Invisible Extreme- ultraviolet Wave http://adsabs.harvard.edu/abs/2014ApJ...786..151S 17

  8. STCE: http://www.stce.be/news/316/welcome.html STCE: http://www.stce.be/news/331/welcome.html STCE: http://www.stce.be/news/274/welcome.html The M2-event finished with an arcade, which is the technical term for a series of post-flare coronal loops. Interestingly, these post-flare loops continued to grow, first reaching the limit of AIA's Field-Of- View (FOV) on 15 October around 17:00UT, then continuing to grow even beyond AIA's FOV. Fortunately, PROBA2's wider-field SWAP telescope came to the rescue and was able to monitor this arcade in its full glory till its disappearance around noon on 17 October. So, the loops of this long duration arcade were visible for about 2.5 days (60 hours!), and at their maximum height, they were towering at least 340.000 km above the solar surface. That's not far from the average Earth-Moon distance! 18

  9. An arcade is a series of post-eruption coronal loops Bastille Day event SOHO: https://soho.nascom.nasa.gov/gallery/Movies/flares.html TRACE: http://soi.stanford.edu/results/SolPhys200/Schrijver/TRACEpodarchive3.html Yashiro et al. (2013): Post-Eruption Arcades and Interplanetary Coronal Mass Ejections http://adsabs.harvard.edu/abs/2013SoPh..284....5Y https://cdaw.gsfc.nasa.gov/publications/yashiro/yashiro2013SolPhys.pdf Two-ribbon flares are characterized by a pair of bright ribbons observed in H-alpha and ultraviolet (UV) images. The ribbons are located on either side of a magnetic polarity inversion line and they separate from each other as the flare progresses. Two-ribbon flares are often associated with filament eruptions and coronal mass ejections (CMEs). After the launch of the filament, long-lived arcades are formed connecting the two ribbons across the polarity inversion line. The emerged assembly of arches is called a post-eruption arcade (PEA). The PEAs are observed at multiple wavelengths and are known also as long-duration (or decay) events (LDEs; Pallavicini, Serio, and Vaiana, 1977) in X-ray observations. The erupting filament becomes the core of the associated CME (Webb and Hundhausen, 1987; Gopalswamy et al. , 2003), thus PEAs are considered as surface signatures of CMEs (Tripathi, Bothmer, and Cremades, 2004). 19

  10. STCE: http://www.stce.be/news/268/welcome.html (concerns the 10 September 2014 flare) One can see that the EUV emissions peak 6-12 minutes later than those from x-ray. This is due to the cooling of the post-flare coronal loops, whose emissions become then better visible in the less energetic EUV passbands. The AIA 094 emissions also show a second peak about 30 minutes after its maximum. This second peak is not visible in x-ray. This "extra" EUV emission does not originate from the original flare site, but most probably from a volume of higher coronal loops. This may indicate there's additional post-flare loop reconnection, but at a lower temperature than during the flare's main peak. This is called the "EUV late phase". Sun et al. (2013): Hot Spine Loops and the Nature of a Late-phase Solar Flare (graph - This concerns the 15 November 2011 flare) http://adsabs.harvard.edu/abs/2013ApJ...778..139S This event also features an extreme-ultraviolet (EUV) late phase, i.e., a delayed secondary emission peak in warm EUV lines (about 2 – 7 MK). We show that this peak comes from the cooling of large post-reconnection loops beside and above the compact fan, a direct product of eruption in such topological settings. The long cooling time of the large arcades contributes to the long delay; additional heating may also be required . … If the fan is small compared to the pre-existing AR, the post-reconnection loops can be very different in size. They will cool at different rates during the initial, conduction-dominated stage, when the cooling time scales with the loop length squared (for recent review, see Reale 2010). Because emission in the warm EUV lines increases only after the hot loops cool down, peaks fromA2 and A3 loops will appear at a much later time compared to A1, as already noted in Woods et al. (2011). Additional heating from ongoing, weak reconnection may also contribute (Hock et al. 2012). The two mechanisms need not be mutually exclusive. NASA: https://svs.gsfc.nasa.gov/10817 The solar EUV radiation creates our Earth's ionosphere (plasma in our atmosphere), so solar flares disturb our ionosphere and consequently our communication and navigation technologies, such as Global Positioning System (GPS), that transmit through the ionosphere. … With these new SDO EVE results, they now recognize that additional ionospheric disturbances from these later EUV enhancements are also a concern. 20

  11. JHV: Flare from 06 June 2000 (SOHO, Yohkoh: X1.1 in NOAA 9026) More on this and other cusps: http://solar.physics.montana.edu/takeda/evt_archive/cusp_flare.html Another example of a cusp: http://www.stce.be/news/298/welcome.html (06 March 2015) Another example of a cusp: http://www.stce.be/news/173/welcome.html (19 January 2012) Another example of a cusp: http://www.stce.be/news/238/welcome.html (25 February 2014; X4.9) Yokoyama et al. (2001): Clear Evidence of Reconnection Inflow of a Solar Flare http://adsabs.harvard.edu/abs/2001ApJ...546L..69Y Magnetic reconnection (Petschek 1964; Sweet 1958; Parker 1963) — the reorganization caused by local diffusion of antiparallel magnetic field lines and the consequent release of magnetic energy — has been thought to be the cause of solar flares (e.g., Shibata 1996). Many indirect pieces of evidence for this process have been found by recent spacecraft observations. There was, however, almost no direct evidence, such as inflow or outflow (reconnection jet) that carries the field lines toward or from the magnetic neutral point where the local dissipation occurs (except for McKenzie & Hudson 1999). We report here the first discovery of reconnection inflow during a flare on 1999 March 18. Solar flares are now thought to be caused by magnetic reconnection (Fig. 1; e.g., Shibata 1996; Yokoyama & Shibata 1998). In this model, the antiparallel field lines dissipate in a certain local point in the corona. The tension force of the reconnected field lines then accelerates the plasma out of the dissipation point. In response to this outflow, the ambient plasma is drawn in. The inflowing plasma carries the ambient magnetic field lines into the dissipating point. These field lines continue the reconnection cycle. In this manner, the magnetic energy stored near the neutral point is released to become the thermal and bulk-flow energy of plasma . … The supporting evidence for this model is the observation of a cusp-shaped soft X-ray flare loop (Tsuneta et al. 1992). The tip of the cusp is thought to be the remnant of the kink of the reconnected field lines. This cusp-shaped flare loop increases its height and the distance between the footpoints, which might be the consequence of the piling up of the shrunk magnetic field lines (see also Forbes&Acton 1996; Hiei, Hundhausen, & Sime 1993). Flare observation by Masuda et al. (1994) demonstrates a hard X-ray source above the soft X-ray loop. This source suggests that some high-energy process, such as acceleration of electrons associated with reconnection, is occurring above the soft X-ray loop. 21

  12. Seaton et al. (2017): Observations of the Formation, Development, and Structure of a Current Sheet in an Eruptive Solar Flare http://adsabs.harvard.edu/abs/2017ApJ...835..139S Figure 2. Evolution of current sheet structure in the 131 Å AIA channel, beginning with its appearance in the wake of a strong CME. Early on (upper left) the structure is long and narrow, and only later (upper right) do background features begin to appear. These features are first seen as shrinking loops, which later broaden (middle left, black arrow) into a more fan-like structure, while the sheet itself (middle left, white arrow) begins to broaden. As the sheet broadens, shrinking loops are clearly visible in the cusp region at the current sheet's base (middle right). Even later (lower left) dark inflows, presumably SADs, become visible in the diffuse background emission. At very late times (lower right) some material appears to flow into the sheet itself, triggering bifurcated up-down flows along the sheet structure. The structures we report on in this paper were formed in association with a large and complex filament eruption that occurred on the east limb of the Sun in NOAA Active Region 11990 at about 00:40 UT on 2014 February 25. This eruption was also associated with an … X4.9 class flare, which peaked about 10 minutes after the onset of the eruption, and a very energetic CME with a reported velocity of more than 2100 km s −1 in the CDAW CME Catalog (for a description of CDAW, see Yashiro et al. 2004). 22

  13. http://www.stce.be/news/157/welcome.html http://www.stce.be/news/218/welcome.html 23

  14. http://www.stce.be/news/255/welcome.html The eruption was not associated to an obvious x-ray flare, but a disturbance was noted in the EUV imagery, parallel to the original position of the erupted filament on both the east and west side (see annotated image above). The disturbance propagated through the corona at a speed of 2-5 km/s. Just as the expanding flare ribbons ("parallel ribbons") and the post-flare coronal loops that often can be seen after a solar flare, also this phenomenon is an effect of the reconnection higher up in the solar atmosphere. The charged particles get accelerated towards the denser inner solar atmosphere, where they collide with other particles and heat the local chromospheric environment and make it evaporate. It is not an EIT wave, characteristics of which were described in a previous newsletter (see http://www.stce.be/news/241/welcome.html for more details). The footpoints of some faint coronal loops can be seen embedded in the expanding disturbance in the combo movies. Another example of a canyon-of-fire NASA: https://www.nasa.gov/content/goddard/nasa-releases-movie-of-suns-canyon-of-fire NASA: https://www.nasa.gov/content/solar-filament-eruption-canyon-of-fire 24

  15. Another name for this kind of features is « solar tsunami » They are expanding large-scale waves in the solar atmosphere usually associated to strong solar eruptions (flares and CMEs). Animation sources: - Moreton wave: 6 December 2006 event: https://en.wikipedia.org/wiki/Moreton_wave#/media/File:MoretonWaveAnimation200612.gif With associated press release at http://www.nso.edu/sites/www.dev.nso.edu/files/files/press/archive/SolarTsunami.pdf - EIT wave: 12 May 1997 event: http://umbra.nascom.nasa.gov/eit/cme/may12/ References: - “EIT waves” and coronal mass ejections, Chen et al. (2011): http://www.ncra.tifr.res.in:8081/~basi/ASICS_2/229-chen.pdf - Synthesis of CME-associated and EIT-wave features from MHD simulations, Chen et al. (2005): http://astronomy.nju.edu.cn/~chenpf/paper/ssr01.pdf - Large-scale coronal propagating fronts… , Nitta et al. (2013) : http://iopscience.iop.org/article/10.1088/0004-637X/776/1/58/pdf - Observation of a Moreton wave and wave- filament interactions… , Liu et al. (2013) : http://iopscience.iop.org/article/10.1088/0004-637X/773/2/166/pdf - On the nature of EIT waves, EUV dimmings and their link to CMEs, Zhukov et al. (2004) : http://www.aanda.org/articles/aa/pdf/2004/44/aa0351-04.pdf - SOHO/EIT Observations of the 1997 April 7 Coronal Transient: Possible Evidence of Coronal Moreton Waves, Thompson et al. (1999): http://adsabs.harvard.edu/abs/1999ApJ...517L.151T 25

  16. More examples: Moreton waves : http://www.stce.be/news/222/welcome.html EIT waves : http://www.stce.be/news/222/welcome.html and http://www.stce.be/news/241/welcome.html EIT-waves are named after the Extreme-ultraviolet Imaging Telescope (EIT) onboard SOHO, with which this phenomenon was discovered in 1996-1997. They are large-scale bright fronts observed in extreme ultraviolet (EUV) and propagating over a significant portion of the solar surface. 17 years later, the true nature of these waves remains under debate, though there is a gradual convergence towards it being primarily a fast magnetosonic wave (directly related to the presence of a coronal mass ejection, CME, rather than a flare), but often also with a contribution from the CME expansion (see Note 1). Other typical characteristics are its relatively low average speed of 200-600km/s, and that these fronts can be stopped at the boundary of coronal holes or near active regions. Note 1 - A fast magnetosonic wave is a longitudinal wave of charged particles in a magnetized plasma propagating in all directions, including perpendicularly and parallel to the magnetic field. See image underneath (Source: Wikipedia). Example above on EIT wave from: https://cor1.gsfc.nasa.gov/ November 24, 2009: Sometimes you really can believe your eyes. That's what NASA's STEREO (Solar Terrestrial Relations Observatory) spacecraft are telling researchers about a controversial phenomenon on the sun known as the "solar tsunami." Years ago, when solar physicists first witnessed a towering wave of hot plasma racing along the sun's surface, they doubted their senses. The scale of the thing was staggering. It rose up higher than Earth itself and rippled out from a central point in a circular pattern millions of kilometers in circumference. Skeptical observers suggested it might be a shadow of some kind — a trick of the eye — but surely not a real wave. "Now we know," says Joe Gurman of the Solar Physics Lab at the Goddard Space Flight Center. "Solar tsunamis are real." The twin STEREO spacecraft confirmed their reality in February 2009 when sunspot 11012 unexpectedly erupted. The blast hurled a billion-ton cloud of gas (a "CME") into space and sent a tsunami racing along the sun's surface. STEREO recorded the wave from two positions separated by 90 degrees, 26

  17. Shen et al. (2014): A Chain of Winking (Oscillating) Filaments Triggered by an Invisible Extreme-ultraviolet Wave http://adsabs.harvard.edu/abs/2014ApJ...786..151S In this paper, we present the observations of a chain of winking filaments and a subsequent jet that are observed right after the X2.1 flare in AR11283. The event also produced an extreme-ultraviolet (EUV) wave that has two components: an upward dome-like wave (850 km s – 1 ) and a lateral surface wave (554 km s – 1 ) that was very weak (or invisible) in imaging observations. By analyzing the temporal and spatial relationships between the oscillating filaments and the EUV waves, we propose that all the winking filaments and the jet were triggered by the weak (or invisible) lateral surface EUV wave. The oscillation of the filaments last for two or three cycles, and their periods, Doppler velocity amplitudes, and damping times are 11-22 minutes, 6-14 km s – 1 , and 25-60 minutes, respectively. We further estimate the radial component magnetic field and the maximum kinetic energy of the filaments, and they are 5-10 G and ~10 19 J, respectively. The estimated maximum kinetic energy is comparable to the minimum energy of ordinary EUV waves, suggesting that EUV waves can efficiently launch filament oscillations on their path. Based on our analysis results, we conclude that the EUV wave is a good agent for triggering and connecting successive but separated solar activities in the solar atmosphere, and it is also important for producing solar sympathetic eruptions. In this paper, we present an interesting observational study of a chain of winking filaments that was in association with a GOES X2.1 flare in the NOAA active region AR11283 (N13W18) on 2011 September 6. The flare was produced with a remarkable EUV wave propagating mainly in the northwest direction, which not only triggered the oscillation of three filaments in the northwest of AR11283, but also launched the oscillation of a long filament and the occurrence of a small jet in the eastern hemisphere, where the wave signature is very weak or even invisible. According to previous studies, winking filaments are often trigged by either chromospheric Moreton waves or coronal EUV waves (Eto et al. 2002; Okamoto et al. 2004). In the present case, we do not detect any significant signature of Moreton waves in Hα observations. On the other hand, as the EUV wave mainly propagated in the northwest of AR11283, it is hard to understand the trigger mechanisms of the F1's oscillation and the occurrence of the small coronal jet. 27

  18. http://www.stce.be/news/362/welcome.html Mason et al. (2016): Relationship of EUV Irradiance Coronal Dimming Slope and Depth to Coronal Mass Ejection Speed and Mass http://adsabs.harvard.edu/abs/2016ApJ...830...20M Large regions of temporary dimming or darkening of preexisting solar coronal emission often accompany coronal mass ejections (CMEs) and may trace field lines opened during the CME. The plasma of the solar corona responds in a number of ways to an eruptive event. Mason et al. (2014) provide details about the physics behind coronal dimming and the observational effects to be considered during analysis. Therein, the case is made for two hypotheses: that the slope of de- convolved, extreme-ultraviolet (EUV) dimming irradiance light curves should be directly proportional to CME speed, and similarly, that dimming depth should scale with CME mass. Dimming regions can be extensive, representing at least part of the “base” of a CME and the mass and magnetic flux transported outward by it. Extensive surveys of EUV images containing coronal dimming events and their relation to CMEs have been performed by Reinard & Biesecker (2008, 2009). For their sample of 100 dimming events, Reinard & Biesecker (2008) found mean lifetimes of 8 hr, with most disappearing within a day. Reinard & Biesecker (2009) studied CMEs with and without associated dimmings, finding that those with dimmings tended to be faster and more energetic. Bewsher et al. (2008) found a 55% association rate of dimming events with CMEs and conversely that 84% of CME events exhibited dimming. The timescale for dimming development is typically several minutes to an hour. This is much faster than the radiative cooling time, which implies that the cause of the decreased emission is more dependent on density decrease than temperature change (Hudson et al. 1996). Studies have demonstrated that dimming regions can be a good indicator of the apparent base of the white light CME (Thompson et al. 2000; Harrison et al. 2003; Zhukov & Auchère 2004). Thus, dimmings are usually interpreted as mass depletions due to the loss or rapid expansion of the overlying corona (Hudson et al. 1998; Harrison & Lyons 2000; Zhukov & Auchère 2004). 28

  19. http://www.stce.be/news/362/welcome.html Cheng et al. (2016): The Nature of CME-flare-Associated Coronal Dimming http://adsabs.harvard.edu/abs/2016ApJ...825...37C During the eruptive events, transient coronal holes, or coronal dimmings, are often observed (Thompson et al. 2000; Harrison et al. 2003; Zhukov & Auchère 2004). Coronal dimming was first observed in Skylab data and characterized as transient coronal holes (Rust & Hildner 1976; Rust 1983). Subsequently, similar observations have been analyzed to study the relationship of dimming with CMEs, flares, and other associated phenomena … By these series of studies, dimmings are usually interpreted as mass depletion due to the loss or rapid expansion of the overlying corona (Hudson et al. 1998; Harrison & Lyons 2000; Zhukov & Auchère 2004). This interpretation is supported by imaging observations of simultaneous and co-spatial dimming in several coronal lines (e.g., Zarro et al. 1999; Sterling et al. 2000), as well as spectroscopic observations (Harra & Sterling 2001; Tian et al. 2012). Although CMEs are also observed to occur without dimming, Reinard & Biesecker (2009) found that non-dimming CMEs all have speeds of less than 800 km s−1, suggesting a more intimate connection between fast CMEs and dimming properties. Krista & Reinard (2013) found further correlations between the magnitudes of dimmings and flares, and the CME mass by studying variations between recurring eruptions and dimmings. Coronal dimming can be produced by various processes, although the main contributor is mass depletion. As summarized by Mason et al. (2014), several different mechanisms have been proposed to explain coronal dimming. (1) Mass-loss dimming: the mass-loss dimming is produced by the ejection of emitting plasma (Harrison & Lyons 2000; Harra & Sterling 2001), which causes darkening of the areas in and near the erupting active region. (2) Thermal Dimming: … this is due to the heating of coronal plasmas to higher temperatures, so that heated areas appear dark in extreme ultraviolet… (3) Obscuration Dimming: … (4) Wave Dimming: … (5) Doppler Dimming: … Among these mechanisms the mass-loss dimming is considered to be the main process of coronal dimming, and this scenario is supported by many recent studies (Sterling & Hudson 1997; Reinard & Biesecker 2008, 2009; Aschwanden et al. 2009). 29

  20. Sterling et al. (2016): Minifilament Eruptions that Drive Coronal Jets in a Solar Active Region http://adsabs.harvard.edu/abs/2016ApJ...821..100S Figure 1. Schematic showing in 2D the formation process of jets, as suggested by Sterling et al. (2015). The bold horizontal black line is the photosphere, curved black lines represent magnetic field that has not undergone magnetic reconnection, curved red lines show field that has undergone reconnection, and red crosses show locations where reconnection is taking place. (a) A compact bipole carrying a mini-filament (blue) resides next to a larger-scale bipole, in a background ambient open coronal field. (b) Due to an unspecified process, the mini- filament-carrying field erupts outward. Its field orientation is such that magnetic reconnection with the surrounding field external to the erupting bipole does not occur, as long as the erupting bipole is on the near side (i.e., the side from which the eruption originated) of the apex of the larger bipole. Reconnection does, however, occur among the field internal to the erupting bipole itself (“internal reconnection”), just as in typical larger scale filament eruptions that result in typical solar flares and CMEs. In the large-scale flares the internal reconnection results in a “normal” solar flare, while in this case the internal reconnection beneath the erupting mini-filament- carrying field results in the jet bright point (JBP) (bold red semicircle). (c) When the erupting field reaches the far end of the larger bipole’s apex, its orientation is favorable for reconnection with the ambient field (“external reconnection”), resulting in a new open field line, and new field loops over the large bipole. A hot coronal jet occurs on the newly reconnected open field lines. (d) If the external reconnection of the ejected mini-filament- carrying field envelope progresses far enough into that field’s core, then the mini-filament material (blue and light blue), which is in the core, will escape along new open field lines, resulting in a cool component of the coronal jet. Example from ”Leaving on a jet”: http://www.stce.be/news/293/welcome.html SWPC glossary: http://www.swpc.noaa.gov/content/space-weather-glossary Surge : A jet of material from active regions that reaches coronal heights and then either fades or returns into the chromosphere along the trajectory of ascent. Surges typically last 10 to 20 minutes and tend to recur at a rate of approximately 1 per hour. Surges are linear and collimated in form, as if highly directed by magnetic fields. 30

  21. Liu et al. (2016): On the Observation and Simulation of Solar Coronal Twin Jets http://adsabs.harvard.edu/abs/2016ApJ...817..126L Decades have passed since the first observations on solar jets (named as surges in Newton 1934), which are thought to play an important role in solar wind acceleration and coronal heating (e.g., Tsiropoula & Tziotziou 2004; Tian et al. 2014). A generalized definition of solar jets includes the terms of Hα surges (e.g., Canfield et al. 1996; Jibben & Canfield 2004), UV/EUV/X-ray jets (e.g., Schmieder et al. 1988; Patsourakos et al. 2008; Tian et al. 2014; Liu et al. 2015) and spicules (e.g., De Pontieu et al. 2007; Shibata et al. 2007), among which their different names come from different dominant temperatures and sizes. As shown in many previous works (Shibata et al. 1996, as a review), different jets obtain quite different physical characteristics such as length and axial speed, which range from a few to hundreds of megameters and tens to thousands of kilometers per second, respectively. Despite the different properties of different jets, it is believed that they are triggered by a similar mechanism (except type I spicules, De Pontieu et al. 2007). Reconnections between newly emerging twisted loops with pre- existing ambient open fields (e.g., Moreno-Insertis et al. 2008) lead to the heating and initiation of bulks of plasma, which are observed as materials of a jet (Savcheva et al. 2007). Twists transferred from the emerging flux then lead to the rotational motion of jets, as observed and studied widely through observation and simulation (e.g., Xu et al. 1984; Shibata & Uchida 1985; Canfield et al. 1996; Shimojo et al. 2007; Pariat et al. 2010; Liu 2009; Fang et al. 2014; Liu et al. 2014). Example from ”Leaving on a jet”: http://www.stce.be/news/293/welcome.html Also: Cheung et al. (2014): Flux Emergence (Theory) http://adsabs.harvard.edu/abs/2014LRSP...11....3C Chapter 4 and Fig. 53. SWPC glossary: http://www.swpc.noaa.gov/content/space-weather-glossary 31

  22. H-alpha flare classification: Australian SWS: http://www.sws.bom.gov.au/Educational/2/4/2 H-alpha observing: http://users.telenet.be/j.janssens/Halpha/Halfaeng.html#Flares From SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 4) Optical Information (Op): The optical classification and location of an associated flare, observed in Hα. It contains an importance and a Brightness parameter: * Importance is the corrected area of the flare in heliospheric square degrees at maximum brightness , observed in the Hα line (656.3 nm). S - Subflare (area ≤ 2.0 deg.2) 1 - Importance 1 (2.1 ≤ area ≤ 5.1 deg. 2) 2 - Importance 2 (5.2 ≤ area ≤ 12.4 deg. 2) 3 - Importance 3 (12.5 ≤ area ≤ 24.7 deg. 2) 4 - Importance 4 (area > 24.8 deg. 2) * Brightness is the relative maximum brightness of flare in Hα. F – faint ; N – normal ; B – brilliant * Location (ºLat. ºCMD) gives the spherical, heliographic coordinates of the solar flare in Hα as a distance in degrees from the solar equator (heliographic latitude), and distance in degrees from a line extending from the north solar rotational pole to the south solar rotational pole through the center of the solar disk as viewed from Earth (central meridian). The field is blank for x-ray events with no optical correlation (no optical flare observed or no optical patrol at the time) and for flares that occasionally occur in unassigned regions). This classification is still widely used, e.g. in the daily SWPC (event) reports, The Weekly, the SIDC’s Ursigrams and weekly bulletins,… A detailed analysis of H-alpha flare properties is by Temmer et al. (2001) Statistical analysis of solar H flares http://www.aanda.org/articles/aa/pdf/2001/33/aa1413.pdf 32

  23. DAY BEGIN MAX END LOC XRAY OP 10CM Catania/NOAA RADIO_BURST_TYPES SF: 22 Apr 2015 - 0830 0844 0858 S09E05 M1.1 SF 1N: 16 Apr 2014 -1954 1959 2004 S14E09 M1.0 1N 24/2035 II/2 2B: 02 Apr 2014 - 1318 1405 1428 N14E53 M6.5 2B 3700 09/2027 II/1IV/2 3B: 07 Mar 2012 - 0002 0024 0040 N17E27 X5.4 3B 7200 IV/1,II/2,V/2 Data are from the SIDC / Daily Ursigrams (http://www.sidc.be/archive ) Images are from GONG/NSO H-alpha Network (ftp://gong2.nso.edu/HA/hag/ ) Cont’s (H-alpha classification) The size (or importance) of a flare can also be measured in MH. From the wikipedia site: https://en.wikipedia.org/wiki/Solar_flare#H-alpha_classification H-alpha classification An earlier flare classification is based on Hα spectral observations. The scheme uses both the intensity and emitting surface. The classification in intensity is qualitative, referring to the flares as: ( f )aint, ( n )ormal or ( b )rilliant. The emitting surface is measured in terms of millionths of the hemisphere and is described below. (The total hemisphere area A H = 15.5 × 10 12 km 2 .) Classification Corrected Area [millionths of hemisphere] S < 100 1 100 – 250 2 250 – 600 3 600 – 1200 4 > 1200 A flare then is classified taking S or a number that represents its size and a letter that represents its peak intensity, v.g.: Sn is a normal subflare . [10] Tandberg-Hanssen, Einar; Emslie, A. Gordon (1988). Cambridge University Press, ed. "The physics of solar flares". 33

  24. Cont’d (H-alpha classification) From Townsend et al. (1982): A source book of the solar-geophysical environment http://www.dtic.mil/dtic/tr/fulltext/u2/a138682.pdf (pp. 105 and 107): One optical flare intensity or "brilliance" classification is based on the Doppler shift of the hydrogen- alpha line. This Doppler shift is a measure of emitting gas particle velocity and is used by the observer in making his subjective estimate of flare intensity. Using this system we classify flares as follows: Intensity Doppler Shift of Flare Emission Faint (F) Seen over a line width of 0.8 Angstrom or greater. Normal (N) Seen over a line width of 1.2 Angstrom or greater. Brilliant (B) Seen at + and/or - 1.0 Angstrom off line center. The SOON (Solar Observing Optical Network) telescopes are capable of directly measuring the intensity of optical flare emissions. The SOON observatories report as their flare brightness the measured flare intensity. However, the observed intensity is strongly dependent on the seeing conditions, and only a slight amount of atmospheric pollution can drastically alter the measured intensity. History of H-alpha observations: http://adsabs.harvard.edu/abs/1966SSRv....5..388S Optical Observations of Solar Flares , Švestka , Zdeněk , 1966 The H-alpha flare classification system was approved by Commission 10 of the IAU in 1966 (Zirin, 1988: Astrophysics of the Sun, pp. 347). 34

  25. Source: http://iopscience.iop.org/article/10.1086/304521/fulltext/36016.text.html From SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 2) The letter classification of solar flares used in these definitions (Table 1) was initiated on 01 January 1969. This classification ranks solar activity by its peak x-ray intensity in the 0.1-0.8 nm band as measured by the Geostationary Operational Environmental Satellites (GOES). This x-ray classification offers at least two distinct advantages compared with the standard optical classifications: it gives a better measure of the geophysical significance of a solar event, and it provides an objective means of classifying geophysically significant activity regardless of its location on the solar disk. Table 1. The SWPC x-ray flare classification Peak Flux Range (0.1-0.8 nm) Classification mks system (W m-2) cgs system (erg cm-2s-1) A Φ <10 -7 Φ <10 -4 B 10- 7 ≤ Φ <10 -6 10- 4 ≤ Φ <10 -3 C 10- 6 ≤ Φ <10 -5 10- 3 ≤ Φ <10 -2 M 10- 5 ≤ Φ <10 -4 10- 2 ≤ Φ <10 -1 X 10- 4 ≤ Φ 10- 1 ≤ Φ The letter designates the order of magnitude of the peak value and the number following the letter is the multiplicative factor. A C3.2 event for example, indicates an x-ray burst with 3.2x10-6Wm-2 peak flux. Solar flare forecasts are usually issued only in terms of the broad C, M, and X categories. Since x-ray bursts are observed as a full-Sun value, bursts below the x-ray background level are not discernible. The background drops to class A level during solar minimum; only bursts that exceed B1.0 are classified as x-ray events. During solar maximum the background is often at the class M level, therefore class A, B, or C x-ray bursts cannot be discerned. Data are measured by the NOAA GOES satellites, monitored in real time in Boulder (Grubb 1975). ------- The C is often referred to as « Common » , M as « Medium (or moderate) », and X as « eXtreme 35

  26. Source: http://iopscience.iop.org/article/10.1086/304521/fulltext/36016.text.html Table 1. The SWPC x-ray flare classification Peak Flux Range (0.1-0.8 nm) Classification mks system (W m-2) cgs system (erg cm-2s-1) A Φ <10 -7 Φ <10 -4 B 10- 7 ≤ Φ <10 -6 10- 4 ≤ Φ <10 -3 C 10- 6 ≤ Φ <10 -5 10- 3 ≤ Φ <10 -2 M 10- 5 ≤ Φ <10 -4 10- 2 ≤ Φ <10 -1 X 10- 4 ≤ Φ 10- 1 ≤ Φ The letter designates the order of magnitude of the peak value and the number following the letter is the multiplicative factor. A C3.2 event for example, indicates an x-ray burst with 3.2x10-6Wm-2 peak flux. Solar flare forecasts are usually issued only in terms of the broad C, M, and X categories. Since x- ray bursts are observed as a full-Sun value, bursts below the x-ray background level are not discernible. The background drops to class A level during solar minimum; only bursts that exceed B1.0 are classified as x-ray events. During solar maximum the background is often at the class M level, therefore class A, B, or C x-ray bursts cannot be discerned. Data are measured by the NOAA GOES satellites, monitored in real time in Boulder (Grubb 1975). X-ray Background: The daily average background x-ray flux as measured by the GOES satellite. To better reflect mid day values, the average is the lower of (a) the average of 1-minute data between 0800UT to 1600UT, or (b) the average of the 0000UT to 0800UT and the 1600UT to 2400UT data. The value is given in terms of x-ray class (Donnelly 1982); (Bouwer, et al.1982). X-ray flux values below the B1 level can be erroneous because of energetic electron contamination of the x-ray sensors. At times of high electron flux at geosynchronous altitude, the x-ray measurements in the low A-class range can be in error by 20-30 percent. Measurements taken during periods of low energetic electron fluxes are much more accurate. 36

  27. Source: http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf Solar Activity in SC24 (Jan 2009 - Dec 2016) Very Low 1291 Low 1214 Moderate 299 High 118 Very High 0 37

  28. SDO/EVE: http://lasp.colorado.edu/eve/data_access/sdo_xray_proxy/eve_goes_xray_proxy PROBA2/LYRA: http://proba2.oma.be/ssa These measure the solar EUV output which is then scaled to GOES so that they can be reliably compared and substituted. So, these scaled EUV measurements are proxies for the GOES x-ray measurements. 38

  29. From the SWPC webpage: NOAA Space Weather Scales The NOAA Space Weather Scales were introduced as a way to communicate to the general public the current and future space weather conditions and their possible effects on people and systems. Many of the SWPC products describe the space environment, but few have described the effects that can be experienced as the result of environmental disturbances. These scales are useful to users of our products and those who are interested in space weather effects. The scales describe the environmental disturbances for three event types: geomagnetic storms, solar radiation storms, and radio blackouts. The scales have numbered levels, analogous to hurricanes, tornadoes, and earthquakes that convey severity. They list possible effects at each level. They also show how often such events happen, and give a measure of the intensity of the physical causes. The « R » stands for Radio Blackout. Note it starts only from M1 class flares and higher. More at http://www.stce.be/news/366/welcome.html 39

  30. From the SWPC webpage: NOAA Space Weather Scales The NOAA Space Weather Scales were introduced as a way to communicate to the general public the current and future space weather conditions and their possible effects on people and systems. Many of the SWPC products describe the space environment, but few have described the effects that can be experienced as the result of environmental disturbances. These scales are useful to users of our products and those who are interested in space weather effects. The scales describe the environmental disturbances for three event types: geomagnetic storms, solar radiation storms, and radio blackouts. The scales have numbered levels, analogous to hurricanes, tornadoes, and earthquakes that convey severity. They list possible effects at each level. They also show how often such events happen, and give a measure of the intensity of the physical causes. The « R » stands for Radio Blackout. Note it starts only from M1 class flares and higher. More at http://www.stce.be/news/366/welcome.html Systematic satellite observations of the Sun started in 1976 with GOES. For each year and for each disturbance type, one can count for every level the number of events. E.g. so far for 2016, we've had only 10 R1 events (flares with intensity between M1 and M5) and 4 R2 events (intensity between M5 and X1). Data can be retrieved at resp. NGDC/NOAA, NASA/NOAA and WDC Kyoto, and run through mid-October 2016. Each graph shows the yearly accumulation of the events, with the yearly International Sunspot Number (SILSO) superposed on it as the gray dashed line. E.g. in the chart above, for 2014 -the year of SC24 maximum-, the number of radio blackouts amounted to 222, consisting of 183 minor (R1), 23 moderate (R2), and 16 strong (R3) events. This is clearly less than during previous solar cycles such as e.g. in 1989 when there were no less than 679 radio blackouts including 59 strong or more intense events! Also, SC24 has not produced any severe or extreme event so far, i.e. X10 or stronger flare. 40

  31. The above chart shows for each bin of solar flare intensity (C3-X6) the ratio of the number of flares for that bin vs. the total number of flares (25031 flares; January 1976 – May 2016). Both axis are logarithmic in nature. The C2 and lower classes were omitted as these numbers are affected during high solar activity (high x-ray background). The X7 and higher intensities were omitted for not sufficient data. The linear expression between these two quantities is y=0.6274-2.1333x This means that the number of flares N for a bin can be calculated from a power law equation: N = 25031. delta . 4.24 I -2.13 , with delta equalling 1, 10 or 100 for the resp. class C, M or X. Another rule of thumb: Since 1976, there have been a total of 55000 x-ray flares. About 48000 were C- class flares, 6500 were M-class flares, and 500 were X-class flares. Or in percentages: For every 100 solar flares, there are 87 C-class flares, 12 M-class flares, and 1 X-class flare. More on this (for the period 1976-1993) is at the Australian SWS: http://www.sws.bom.gov.au/Educational/2/4/5 41

  32. From SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 15) The start of an x-ray event is defined as the first minute in a sequence of 4 minutes of steep monotonic increase in 0.1-0.8 nm flux. The time of x-ray maximum is defined as the time tag of the peak 1-minute averaged value x-ray flux. The end time is the time when the flux level decays to a point halfway (1/2 peak) between the maximum flux and the pre-flare background level. 42

  33. From SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 15) The start of an x-ray event is defined as the first minute in a sequence of 4 minutes of steep monotonic increase in 0.1-0.8 nm flux. The time of x-ray maximum is defined as the time tag of the peak 1-minute averaged value x-ray flux. The end time is the time when the flux level decays to a point halfway (1/2 peak) between the maximum flux and the pre-flare background level. From Temporal aspects and frequency distributions of solar soft X-ray flares Veronig et al. (2002): http://www.aanda.org/articles/aa/pdf/2002/06/aa1910.pdf And from The duration of solar flares http://www.stce.be/news/332/welcome.html 43

  34. From the SWPC glossary at http://www.swpc.noaa.gov/content/space-weather-glossary#longduration (operational definition) long duration event (LDE) With reference to x-ray events, those events that are not impulsive in appearance. The exact time threshold separating impulsive from long-duration events is not well defined, but operationally, any event requiring 30 minutes or more to decay to one-half peak flux is regarded as an LDE. It has been shown that the likelihood of a coronal mass ejection increases with the duration of an x-ray event, and becomes virtually certain for durations of 6 hours or more. 44

  35. From the SWPC glossary at http://www.swpc.noaa.gov/content/space-weather-glossary#longduration (operational definition) long duration event (LDE) With reference to x-ray events, those events that are not impulsive in appearance. The exact time threshold separating impulsive from long-duration events is not well defined, but operationally, any event requiring 30 minutes or more to decay to one-half peak flux is regarded as an LDE. It has been shown that the likelihood of a coronal mass ejection increases with the duration of an x-ray event, and becomes virtually certain for durations of 6 hours or more. Imagery from STCE: http://www.stce.be/news/332/welcome.html A short and a long duration X1 flaring event. These took place resp. in NOAA 11890 on 10 November 2013 (duration: 10 minutes) and in NOAA 11520 on 12 July 2012 (duration: 113 minutes or nearly 2 hours). The latter was accompanied by a full halo CME (no surprize), but also the 2013 X1 flare was associated with a partial halo CME. 45

  36. From SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 1) Terms Used to Describe Solar Activity Very Low: x-ray events less than C-class. Low: C-class x-ray events. Moderate: isolated (one to four) M-class x-ray events. High: several (5 or more) M-class x-ray events, or isolated (one to four) M5 or greater x-ray events. Very High: several (5 or more) M5 or greater x-ray events. Wheatland et al. (2002): Understanding solar flare waiting-time distributions http://www.physics.usyd.edu.au/wheat/papers/pdfs/understanding_WTD.pdf Figure 2 shows the waiting-time distributions for the GOES events (greater than C1 class) for the maximum and minimum phases of the cycle as defined by Figure 1. The upper panel shows the WTD for all years 1975 – 2001, and reproduces the power-law tail reported by Boffeta et al. (1999). The dashed vertical line indicates the average waiting time, which is about 6.5 hours. The lower panel shows the WTDs for the maximum and minimum phases of the cycle. The distribution for the maximum phase has a steeper distribution, because the rate of flaring is higher around solar maximum, and so the average waiting time is less. The average waiting times for the two phases are indicated by the dashed vertical lines. The maximum and minimum distributions both exhibit approximate power-law tails. From http://users.telenet.be/j.janssens/Archives/Archives.html#021109 The longest stretch without C-class flares was from 3 April till 3 November 2008, that’s 214 consecutive days of very low activity. Since the start of systematic GOES observations, there have been only 9 periods with more than 60 consecutive days with no C-class flares, 6 of those happened during the most recent SC23-SC24 minimum … The longest stretch without M-class flares was from 25 March 2008 till 19 January 2010 (665 days ). … The longest stretch without X- class flares was from 14 Dec 2006 till 15 February 2011 (1524 days). 46

  37. More info at http://www.stce.be/news/244/welcome.html And in Ranns et al., 2000, Emerging flux as a driver for homologous flares http://adsabs.harvard.edu/abs/2000A%26A...360.1163R And in Zirin, Astrophysics of the Sun, 1988 pp. 353 Homologous flares are the solar equivalent of identical twins. They concern a series of solar flares taking place repetitively in the same active region with essentially the same position and with a common pattern of development, i.e. having the same main footpoints and general shape in the main phase as defined in H-alpha or EUV-imagery. Though not a requirement, homologous flares often have similar strength, and if there are more than two, they sometimes occur within similar time intervals. Image source: Solar Flares, Zdenek Svestka, 1976, chapter II, page 24 47

  38. More at http://www.stce.be/news/249/welcome.html Last week, scientists got a few additional simultaneous flares requiring further investigation. No less than 8 flare events had coinciding brightenings in two or even three well separated sunspot regions. Half of these occurred between NOAA 2035 and 2038. This movie shows two examples on 22 and 23 April. It concerns a C2 and C4 flare peaking resp. at 18:41UT and 01:04UT (images underneath). In both cases, the brightening peaked almost at the same moment in the sunspot groups. Also Moon et al. (2002): Statistical Evidence for Sympathetic Flares http://iopscience.iop.org/article/10.1086/340945/fulltext/55477.text.html Global connection: 1 August 2010 event: http://science.nasa.gov/science-news/science-at-nasa/2010/13dec_globaleruption/ 48

  39. Imagery from STCE: http://www.stce.be/news/218/welcome.html In contrast to the eruption 5 days earlier, the 29 September event resulted in a minor C-class x-ray flare (a so- called Hyder flare, albeit a weak one). It concerned a long duration event that started at 21:43UT, so about 15 minutes after the first visible signs of the eruption in H-alpha and EUV. The flare reached its maximum at 23:39UT and lasted 200 minutes (over 3 hours!). Numerous post-flare coronal loops were visible. Another difference was that this eruption was also associated to a moderate proton event. This was only the fifth such event this year. Info at STCE: http://www.stce.be/news/281/welcome.html And at Australian SWS: http://www.sws.bom.gov.au/Educational/2/4/1 Another example at http://www.stce.be/news/203/welcome.html Two seminal papers by C. Hyder (theory not entirely correct!) A Phenomenological Model for Disparitions Brusques followed by Flarelike Chromospheric Brightenings, I: The Model, its Consequences, and Observations in Quiet Solar Regions http://adsabs.harvard.edu/abs/1967SoPh....2...49H (1967) A Phenomenological Model for disparitions brusques followed by Flarelike Chromospheric Brightenings, II: Observations in Active Regions http://adsabs.harvard.edu/abs/1967SoPh....2..267H (1967) Luo (1982): The flares of spotless regions http://adsabs.harvard.edu/abs/1982AcASn..23...95L The 20 flares of sunspotless regions observed at Yunnan Observatory in cycle 21 are analyzed. It is found that the natural productivity of spotless flares is about three percent, that the distribution of the Carrington longitudes tends to shift eastwards, that most of the spotless region flares are those at low energy, and that the background conditions of producing spotless region flares are the same as those of producing spot region flares. Thus, there must be a magnetic field structure with opposite polarity in the solar atmosphere of a flare region. The change in inclination of a filament to a fibril in a spotless region from a large angle to a small one indicates that the force exerted on the spotless active region transforms gradually from pressure to shearing force, meaning that the activity changes gradually from energy storage to sudden energy release. 49

  40. NASA: Hinode Discovers the Origin of White Light Flare (2010) https://www.nasa.gov/centers/marshall/news/news/releases/2010/10-052.html White light emissions were observed by the Solar Optical Telescope during an X-class flare that occurred at 22:09 UT on Dec. 14, 2006 (see Fig. 1). The RHESSI satellite simultaneously recorded hard X-ray emissions, an indicator of non-thermal electrons accelerated by solar flares. The team found that the spatial location and temporal change of white light emissions are correlated with those of hard X- ray emissions (see Fig. 2). Moreover, the energy of white light emissions is equivalent to the energy supplied by all the electrons accelerated to above 40 keV (~40 percent of the light speed). This finding strongly suggests that highly accelerated electrons are responsible for producing white light emissions. Hard X-rays are emitted when accelerated electrons impact the dense atmosphere near the solar surface. Normally, white light emissions primarily come from the solar surface, whereas 40 keV electrons can penetrate into the atmosphere about 1,000 km above the solar surface, i.e., the chromosphere. Fig. 1, above, White light images of solar surface observed by the Hinode Solar Optical Telescope at 22:07 UT, before the flare, and below, at 22:09 UT, during the flare on Dec. 14, 2006. Image Credit: NASA/JAXA Fig.2: White light emission, left, taken by Hinode/SOT, and the difference image of white light emission and RHESSI hard X-ray contours at 22:09 UT. The background image is the differential white light image (the average of the images taken at 22:07 UT and 22:17 UT is subtracted). Blue contours show 40-100 keV emission. Image credit: NASA/JAXA An overview of WLFs is at http://users.telenet.be/j.janssens/WLF/Whitelightflare.html (last update: August 2012). 50

  41. Gamma-ray Burst Effects on the Ionosphere http://vlf.stanford.edu/research/gamma-ray-burst-effects-ionosphere http://news.stanford.edu/news/2006/march1/ainansr-030106.html Inan et al. (2007): Massive disturbance of the daytime lower ionosphere by the giant g-ray flare from magnetar SGR 1806-20 http://adsabs.harvard.edu/abs/2007GeoRL..34.8103I Wiki: The intense radiation of most observed GRBs is believed to be released during a supernova or hypernova as a rapidly rotating, high-mass star collapses to form a neutron star, quark star, or black hole. Gamma-ray flares: NASA's Fermi Detects the Highest-Energy Light from a Solar Flare http://svs.gsfc.nasa.gov/vis/a010000/a011000/a011000/index.html Solar flares produce gamma rays by several processes, one of which is illustrated here. The energy released in a solar flare rapidly accelerates charged particles. When a high-energy proton strikes matter in the sun's atmosphere and visible surface, the result may be a short-lived particle — a pion — that emits gamma rays when it decays. This image from Fermi's Large Area Telescope (LAT) shows how the entire sky looked on March 7 in the light of gamma rays with energies beyond 100 MeV. Although the Vela pulsar is the brightest continuous LAT source, it was outmatched this day by the X5.4 solar flare, which brightened the sun by 1,000 times. Nuclear radiation: http://www.world-nuclear.org/information-library/safety-and-security/radiation-and- health/nuclear-radiation-and-health-effects.aspx Nuclear radiation arises from hundreds of different kinds of unstable atoms. While many exist in nature, the majority are created in nuclear reactions. Ionizing radiation which can damage living tissue is emitted as the unstable atoms (radionuclides) change ('decay') spontaneously to become different kinds of atoms. Gamma ray event lists and plots can be found at https://hesperia.gsfc.nasa.gov/fermi_solar/ 51

  42. Ajello et al. (2014): Impulsive and Long Duration High-energy Gamma-Ray Emission from the Very Bright 2012 March 7 Solar Flares http://adsabs.harvard.edu/abs/2014ApJ...789...20A Another example of a gamma flare (20 January 2005) https://www.nasa.gov/vision/universe/solarsystem/solar_fireworks.html https://www.nasa.gov/home/hqnews/2005/may/HQ_05132_solar_fireworks.html Movies at https://svs.gsfc.nasa.gov/3162 “This flare produced the largest solar radiation signal on the ground in nearly 50 years," said Dr. Richard Mewaldt of the California Institute of Technology in Pasadena. Normally it takes two or more hours for a dangerous proton shower to reach maximum intensity at Earth after a solar flare, but the particles from the January 20 flare peaked about 15 minutes after the first sign. That's important," said Mewaldt, "because it's too fast to respond with much warning to astronauts or spacecraft that might be outside Earth's protective magnetosphere. In addition to monitoring the Sun, we need to develop the ability to predict flares in advance if we are going to send humans to explore our solar system.“ The event also shakes current theory about the origin of proton storms at Earth. "Since about 1990, we've believed that proton storms at Earth are caused by shock waves in the inner solar system as coronal mass ejections plow through interplanetary space," says Professor Robert Lin of the University of California at Berkeley, principal investigator for the Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI). "But the protons from this event may have come from the Sun itself, which is very confusing." TRACE website: http://www.lmsal.com/TRACE/ ; http://www.lmsal.com/TRACE/POD/TRACEpodarchive21.html#case03 Active region NOAA 10720 is one of the most flare-productive regions of the last few years, with 10 M-class and 5 X-class flares in a week. The largest flare (to date) from the region was an X7.9 on 20-Jan-2005, starting at 06:36 UT. The 53 min. movie (11MB) showing the UV evolution of the flare (C IV 1600Å band). Movie courtesy of Dawn Myers (GSFC). The proton storm associated with this flare impacted many spacecraft. Not only those orbiting Earth were affected: the NASA Deep Impact mission en route to comet Tempel 1, even went into a safehold until the electronics could be restarted after the storm passed. 52

  43. See notes at: STCE news item: http://www.stce.be/news/374/welcome.html SWS: http://www.sws.bom.gov.au/Educational/2/2/5 SWS: http://www.sws.bom.gov.au/Educational/2/2/6 Tapping, K. (2013): The 10.7 cm solar radio flux (F 10.7 ) http://adsabs.harvard.edu/abs/2013SpWea..11..394T The connection between sunspots and solar centimetric emissions was discovered independently by Covington [1947, 1948], Lehaney and Yabsley [1949], and through a statistical study, Denisse [1948]. Covington [1947] used the edge of the Moon during a solar eclipse to identify a significant emission contribution associated with a large active region. The utility of what became known as F10.7 as an indicator of the level of solar activity led to the continuation of measurement to the present day and to the program becoming a data service. A 10.7 cm solar flux measurement is a determination of the strength of solar radio emission in a 100 MHz-wide band centered on 2800 MHz (a wavelength of 10.7 cm), averaged over an hour. It is expressed in solar flux units (sfu), where 1 sfu = 10 – 22 W m – 2 Hz – 1. This data filtering procedure has been discontinued for two reasons: one was the staffing issue mentioned earlier. Second, many applications require the measured flux value, not a value that has been modified. Subsequently, practice has been to distribute the data as measured and to provide auxiliary data so that users could apply whatever data modification procedures they require. Three flux determinations are made each day, at 1700, 2000, and 2300 UT, except during the winter months, where the low elevation of the Sun (DRAO lies at +50° latitude) and the hilly terrain, forces the times to be changed to 1800, 2000, and 2200 UT. Each flux determination takes 1 h and takes no account of the solar radio emissions recorded outside the intervals covered by the measurements. Since the active region emissions contributing to the slowly varying emission (and F10.7)may vary over hours or less, there may be a significant degree of undersampling. In addition, there could be a contribution by a burst. The undersampling means there is a possible error if one uses a flux value in an application involving a different time from that at which the flux measurement is made. 53

  44. 1 sfu = 10-22 Wm-2 Hz-1 Tapping (2013): The 10.7 cm solar radio flux (F10.7) http://onlinelibrary.wiley.com/doi/10.1002/swe.20064/epdf A 10.7 cm solar flux measurement is a determination of the strength of solar radio emission in a 100 MHz-wide band centered on 2800 MHz (a wavelength of 10.7 cm), averaged over an hour. It is expressed in solar flux units (sfu), where 1 sfu = 10 – 22 Wm – 2 Hz – 1. It is daily measured at Penticton, British Columbia, Canada (DRAO: Dominion Radio Astrophysical Observatory). Measurements are taken at 17UT, 20UT and 23UT (winter period: 18-20-22UT), with the local noon value (20UT) as the value for that day. It is uncorrected for any flare influence. The daily values are at http://www.spaceweather.ca/solarflux/sx-4a-en.php The 10.7cm radio flux consists of three identifiable components: a rapidly varying or R component, comprising emissions varying over timescales in the second-minute range, perhaps as long as an hour. Slower variations were lumped into a slowly varying or S component. Extrapolation to zero activity suggested an underlying constant, base level, which became called the quiet sun, or Q component. The terms R and Q have fallen out of use, and these components are now known, respectively, as bursts and the quiet sun background emission. The slowly varying component originates primarily in active regions; its intensity is a measure of the overall level of solar magnetic activity and has a broad spectral peak at about 10 cm wavelength. The F10.7 values comprise contributions from the S component and the quiet sun background, and sometimes from radio bursts. From SWPC Glossary at http://www.swpc.noaa.gov/content/space-weather-glossary#t Tenflare: A solar flare accompanied by a 10cm radio burst of intensity greater than 100% of the pre- burst value. 54

  45. An example of a radio flare (17 May 2013) http://link.springer.com/article/10.1007/s11038-008-9265-8 http://www.spaceweather.org/ISES/code/aaf/ugeoi.html http://www.swpc.noaa.gov/content/space-weather-glossary#t http://www.solarham.net/tenflare.htm 5490 0843 0857 0919 G15 5 XRA 1-8A M3.2 4.4E-02 1748 5490 + 0847 0853 0913 SVI G RBR 4995 800 1748 5490 + 0848 0857 0912 SVI G RBR 2695 450 1748 5490 + 0848 0848 0848 SVI G RBR 410 100 1748 5490 + 0848 0853 0912 SVI G RBR 8800 620 1748 5490 + 0850 //// 1120 SVI C RSP 025-180 IV/2 1748 5490 + 0850 //// 0918 SVI U RSP 025-180 II/2 376 1748 5490 + 0850 0855 0912 SVI G RBR 15400 410 1748 5490 + 0850 0858 0912 SVI G RBR 1415 190 1748 5490 + 0851 0852 0907 SVI G RBR 245 1500 1748 5490 0853 0855 0857 LEA G RBR 610 210 1748 5490 + B0854 U0854 1056 SVI 3 FLA N12E57 2B PRB 1748 5490 B0912 //// A1319 SOH 4 CME XUV,EUV,UV061-060/FS1436 1748 This flare occurred around 09UT, and had little influence on the official radio flux for that day (20UT): Date Time Julian day Carrington rotation Observed Flux Adjusted Flux URSI Flux 2013-05-17 17:00:00 2456430.197 2137.106 137.9 141.1 127.0 2013-05-17 20:00:00 2456430.322 2137.111 136.4 139.6 125.6 2013-05-17 23:00:00 2456430.447 2137.115 135.7 138.8 125.0 Data from RNCan: http://www.spaceweather.ca/solarflux/sx-5-flux-en.php?year=2013 From the SWPC’s “Solar Events” User guide (ftp://ftp.swpc.noaa.gov/pub/indices/events/README ): RBR: The peak value above pre-burst background of associated radio bursts at frequencies 245, 410, 610, 1415, 2695, 4995, 8800 and 15400 MHz: 1 flux unit = 10-22 Wm-2 Hz-1 55

  46. Source of Figure: SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 5) Mind the orientation of the vertical axis! Other figures may have a reversed direction. As the frequency is proportional to the square root of the density, and the density decreases with increasing distance from the Sun, a decreasing frequency means locations higher up in the solar atmosphere. The ionospheric cut-off frequency is around 15MHz (due to too low frequency and so reflected by ionosphere). In order to observe radio disturbances below this frequency, one has to use satellites (above the earth atmosphere) such as STEREO/SWAVES or WIND. Radio bursts at low frequencies (< 15 MHz) are of particular interest because they are associated with energetic CMEs that travel far into the interplanetary (IP) medium and affect Earth’s space environment if Earth-directed. Low frequency radio emission needs to be observed from space because of the ionospheric cutoff. Example: https://stereo-ssc.nascom.nasa.gov/browse/2017/01/16/insitu.shtml Solar Radio Bursts and Space Weather, S.M. White https://www.nrao.edu/astrores/gbsrbs/Pubs/AJP_07.pdf White: Solar radio bursts at frequencies below a few hundred MHz were classified into 5 types in the 1960s (Wild et al., 1963). Coronal Mass Ejections and solar radio emissions, N. Gopalswamy http://citeseerx.ist.psu.edu/viewdoc/download?doi=10.1.1.708.626&rep=rep1&type=pdf Gopalswamy: The three most relevant to space weather radio burst types are type II, III, and IV. Three types of low-frequency non-thermal radio bursts are associated with coronal mass ejections (CMEs): Type III bursts due to accelerated electrons propagating along open magnetic field lines, type II bursts due to electrons accelerated in shocks, and type IV bursts due to electrons trapped in post-eruption arcades behind CMEs. [Radio burst type II, III, and IV are also the only ones that ever get mentioned in the Ursigrams. ] 56

  47. Source - Table taken from the Australian SWS: http://www.sws.bom.gov.au/Category/World%20Data%20Centre/Data%20Display%20and%20Downlo ad/Spectrograph/Solar%20Radio%20Burst%20Classifications.pdf From the SWPC’s “Solar Events” User guide (ftp://ftp.swpc.noaa.gov/pub/indices/events/README ): RSP: Type/Intensity Type II: Slow drift burst Type III: Fast drift burst Type IV: Broadband smooth continuum burst Type V: Brief continuum burst, generally associated with Type III bursts Type VI: Series of Type III bursts over a period of 10 minutes or more, with no period longer than 30 minutes without activity Type VII: Series of Type III and Type V bursts over a period of 10 minutes or more, with no period longer than 30 minutes without activity Type CTM: Broadband, long-lived, decametric continuum Intensity is a relative scale 1=Minor, 2=Significant, 3=Major Shock speed in km/s Note from Dr Christophe Marqué (ROB): Types VI and VII are not used outside NOAA reports. They are not "official" within the radio community. 57

  48. Source: https://www.law.cornell.edu/cfr/text/47/2.101 Question: The 10.7cm radio flux belongs to which metric subdivision? 58

  49. Image courtesy: GOES-curve: STAFF viewer, http://www.staff.oma.be Radio plot: ROB/Humain Radio Observatory, http://www.sidc.be/humain/ 13 June 2014 3940 1521 1524 1527 G15 5 XRA 1-8A C2.4 5.2E-04 2087 3940 + 1521 1522 1523 SAG G RBR 245 290 2087 3940 + 1521 //// 1523 SAG C RSP 025-180 III/2 2087 3940 + 1522 1522 1525 HOL 3 FLA S19E38 SF 2087 Solar Radio Bursts and Space Weather, S.M. White https://www.nrao.edu/astrores/gbsrbs/Pubs/AJP_07.pdf White: Type III bursts are brief radio bursts that drift very rapidly in frequency versus time (Fig. 1). For example, it can drift from 50 to 20 MHz in about 3 seconds, or 10 MHz s−1. Type IIIs are commonly seen in the impulsive phase of solar flares, and the connection they imply between the acceleration region in solar flares and open field lines that reach the solar wind makes them important for understanding field line connectivity in flares and the access of flare – accelerated particles to the Earth. Coronal Mass Ejections and solar radio emissions, N. Gopalswamy http://citeseerx.ist.psu.edu/viewdoc/download?doi=10.1.1.708.626&rep=rep1&type=pdf 59

  50. Culgoora spectrograph at 01 Nov 2003 - http://www.sws.bom.gov.au/Solar/2/2/1 Solar Radio Bursts and Space Weather, S.M. White https://www.nrao.edu/astrores/gbsrbs/Pubs/AJP_07.pdf Type II bursts typically occur at around the time of the soft X – ray peak in a solar flare and are identified by a slow drift to lower frequencies with time in dynamic spectra, the frequent presence of both fundamental and second – harmonic bands (with a frequency ratio of 2), and splitting of each of these bands into two traces. The frequency drift rate is typically two orders of magnitude slower than that of the (“fast–drift”) Type III bursts, so the two burst types are readily distinguished. Hillan et al. (2012): Type II solar radio bursts: Modeling and extraction of shock parameters http://onlinelibrary.wiley.com/doi/10.1029/2011JA016754/full Coronal mass ejections (CME) driving shocks through the corona and into the heliosphere have long been associated with interplanetary type II bursts in the kilometric range [ Wild et al. , 1963; Cane et al. , 1982; Cane , 1985; Nelson and Melrose , 1985; Reiner et al. , 2001]. Blast waves have long been discussed as potential shock drivers in the metric type II burst range [ Cliver et al. , 1999; Claßen and Aurass , 1999], but are not usually thought to persist into the interplanetary medium to drive kilometric emission [ Cane et al. , 1987]. Similarly, it is not clear that all metric type IIs are associated with CMEs, since the metric emission does not routinely (if ever) continue smoothly to the kilometric emission of an interplanetary type II [ Cane and Erickson , 2005]. In the foregoing type II theories, it is the presence and characteristics of the shock that are important, not the mechanism which produced it. [The fundamental band is the one provoked by the shock of the CME and is the one that reaches the lowest frequencies first (track « B » in the image). It is the fundamental track that is used to calculate the (true) speed of the shock as it moves through the corona and away from the Sun (density decrease => frequency decrease).] [The particles from the solar eruption disturb the environment of the particles already present in the higher-up corona. These particles start to oscillate, creating Langmuir waves in the process. These waves generated by both populations of particles, can interact with each other in different ways, as the particles don’t move together and more or less stay at the same place. From these wave interactions, the fundamental and harmonic radio emissions are produced, i.e. at the local plasma frequency and multiples from it.] 60

  51. [Jasmina Magdalenic PhD: Plasma emission is the dominant coherent emission process for the majority of solar radio bursts at decimeter and longer wavelengths. The plasma emission may be defined as any emission process in which the energy of the Langmuir turbulence is partly converted into the escaping radiation. The plasma emission is a multi-stage process, including: - Formation of an unstable beam distribution by the velocity dispersion. - Generation of the Langmuir turbulence as a consequence of plasma instabilities. - Nonlinear evolution and conversion into escaping electromagnetic radiation. Two steps can be distinguished: a) conversion into fundamental transverse radiation - fundamental plasma emission. The fundamental plasma emission at the frequency fp ~ wp/2, is due to conversion into escaping radiation with only small changes in frequency. This conversion could be: 1) scattering of Langmuir waves into transverse waves of thermal ions, 2) coalescence of Langmuir waves and low-frequency waves, such as ion-acoustic waves, into transverse waves, and 3) sort of “direct” conversion due to plasma inhomogeneities. b) production of secondary Langmuir waves and generation of second-harmonic transverse radiation - second harmonic emission. Coalescence of two Langmuir waves results in escaping radiation at the sum of their frequencies (f ~ 2fp), named the harmonic emission. Roberts (1959): Solar Radio Bursts of Spectral Type II : http://adsabs.harvard.edu/abs/1959AuJPh..12..327R Gopalswamy: Coronal Mass Ejections and solar radio emissions : http://citeseerx.ist.psu.edu/viewdoc/download?doi=10.1.1.708.626&rep=rep1&type=pdf 61

  52. Gopalswamy: Coronal Mass Ejections and solar radio emissions http://citeseerx.ist.psu.edu/viewdoc/download?doi=10.1.1.708.626&rep=rep1&type=pdf The type IV bursts are associated with very energetic CMEs (average speed 1200 km/s), confirming the earlier finding by Robinson [1986] for the continuum events at metric wavelengths. The radio emission should originate from a heliocentric distance 3.5 to 4.5 Rs, depending on whether the radio emission occurs at the fundamental or harmonic of the plasma frequency. When the type IV burst attains the lowest frequency, the IP type II burst occurs at frequencies well below 1 MHz, which means the shock is much farther away. This suggests that the energetic electrons responsible for the type IV burst might come from the continued reconnection occurring beneath the CME. [Comment by Dr Christophe Marqué (ROB): The height of type IV reported by Gopalswamy concerns the low frequency ones. The one for example observed in Humain (04 Nov 2015) is really taking place in the post flare loops close to the flare site.] Solar Radio Bursts and Space Weather, S.M. White https://www.nrao.edu/astrores/gbsrbs/Pubs/AJP_07.pdf Type IV bursts are broadband quasi – continuum features associated with the decay phase of solar flares. They are attributed to electrons trapped in closed field lines in the post – flare arcades produced by flares; their presence implies ongoing acceleration somewhere in these arcades, possibly at the tops of the loops in a “helmet–streamer” configuration. Type IV bursts have long been of interest in Space Weather studies because they have a high degree of association with solar energetic particle events. Example: 04 Nov 2015: http://www.stce.be/news/326/welcome.html 2340B1327 U1339 A1348 SVI 2 FLA N09W04 2B ERU 2443 2340 + 1331 1352 1413 G15 5 XRA 1-8A M3.7 5.9E-02 2443 2340 + 1336 1341 1438 SVI G RBR 4995 740 2443 2340 + 1337 1341 1442 SVI G RBR 2695 340 2443 2340 + 1337 1341 1429 SVI G RBR 8800 560 2443 2340 + 1338 1341 1414 SVI G RBR 15400 210 2443 2340 + 1343 //// 1358 SAG C RSP 048-180 II/2 955 2443 2340 + 1351 //// 1531 SVI C RSP 025-171 IV/1 2443 2340 + 1404 1426 1502 SAG G RBR 410 1400 2443 2340 + 1405 1433 1507 SAG G RBR 245 1400 2443 2340 + 1406 1427 1456 SAG G RBR 1415 5800 2443 2340 + 1406 1427 1458 SAG G RBR 610 1000 2443 62

  53. McIntosh, P.S. (1990): The classification of sunspot groups http://adsabs.harvard.edu/abs/1990SoPh..125..251M 63

  54. McIntosh, P.S. (1990): The classification of sunspot groups http://adsabs.harvard.edu/abs/1990SoPh..125..251M Questions to ask (Table 1 from McIntosh paper) Z – General outlook of the sunspot group: => Unipolar or bipolar group? => Penumbra or no penumbra? => Penumbra on one or both sides of the group? => Length of the group (>10°? >15°?) * 7 options: A, B, C, D, E, F, H p – Penumbra largest spot => Rudimentary or mature penumbra? => Symmetric or asymmetric penumbra main spot? => N-S-diameter of the largest spot (>2,5°?) * 6 options: x, r, s, a, h, k c – Sunspot distribution interior (“compactness”) => Several spots between leading and trailing main spot? => Internally, is there at least one spot with a mature penumbra? * 4 options: x, o, i, c (open, intermediate, compact) 64

  55. Bloomfield S. et al. (2012): Toward Reliable Benchmarking of Solar Flare Forecasting Methods http://adsabs.harvard.edu/abs/2012ApJ...747L..41B 65

  56. Bloomfield S. et al. (2012): Toward Reliable Benchmarking of Solar Flare Forecasting Methods http://adsabs.harvard.edu/abs/2012ApJ...747L..41B 66

  57. McIntosh, P.S. (1990): The classification of sunspot groups http://adsabs.harvard.edu/abs/1990SoPh..125..251M 6 July 2012: 8 C, 6 M, X (X1) 67

  58. Jaeggli and Norton (2016): The Magnetic Classification of Solar Active Regions 1992-2015 http://adsabs.harvard.edu/abs/2016ApJ...820L..11J http://iopscience.iop.org/article/10.3847/2041-8205/820/1/L11/pdf Magnetic classifications provide a simple way to describe the configuration of the magnetic flux and sunspots in a solar active region (AR). The Mount Wilson (or Hale) classification system for sunspot groups put forward by Hale et al. (1919) has been used for nearly a century. In the original Hale classification scheme, the designation ( alpha ) is given to regions that contain a single sunspot or sunspot group all having the same polarity. Generally, these also have a weaker opposite polarity counterpart that is not strong or concentrated enough to produce sunspots. ( beta ) is assigned to regions that have two sunspots or sunspot groups of opposite polarity. The classification ( gamma ) is appended to the above classes to indicate the AR has a complex region of sunspots with intermixed polarity. This classification can also be used individually to describe an AR that has no organized magnetic behavior. As an addendum to the original scheme, Kunzel (1965) proposed an additional classification to modify the existing three. ( delta ) indicates that at least one sunspot in the region contains opposite magnetic polarities inside of a common penumbra separated by no more than 2° in heliographic distance (24 Mm or 33″ at disk center). Also at STCE: http://www.stce.be/news/222/welcome.html Make sure to avoid classifying too quickly a sunspot group as a delta or a gamma type when this sunspot group is still very close to the limb. Indeed, line-of-sight may come into play that show an unipolar spot as if it would have a delta structure. See STCE: http://www.stce.be/news/188/welcome.html The pictures to the right are from SDO/HMI and show a magnetogram and a white light image of NOAA 1875 on 23 October 2013. 68

  59. Figures from Townsend et al. (1982): A source book of the solar-geophysical environment http://www.dtic.mil/dtic/tr/fulltext/u2/a138682.pdf See STCE: http://www.stce.be/news/188/welcome.html Text underneath based on SWPC User Guide, SWPC Glossary (http://www.swpc.noaa.gov/content/space-weather-glossary#m ), Mount Wilson (http://obs.astro.ucla.edu/spotlgnd.html ) and SIDC old webpages (http://sidc.oma.be/educational/classification.php#magnetic ). Alpha - Unipolar group; that is, all plus or all minus magnetic field Beta - A bipolar group; that is a mix of plus and minus magnetic polarities exist, with the plus well divided from the minus with one polarity in each end (E- W) of the group, i.e. “easily divided by a simple line”. Beta-Gamma - A group which is generally bipolar but which is lacking a well marked dividing line between the opposite polarity regions (“you need to lift your pencil to divide the polarities” or “ no single, continuous line can be drawn between spots of opposite polarities” ). Gamma - a group in which the polarities are so completely mixed that no bipolar structure can obviously be recognized. Delta - This is a sub-classification for non-unipolar regions. It means at least two opposing polarity umbrae are within two heliographic degrees of each other and share the same penumbra. The chart with the % of Hale classification was based on the NOAA reports at https://solarscience.msfc.nasa.gov/greenwch.shtml for the period Jan 1996-September 2016. A total of 32965 classifications were made. The percentage of reported delta’s in the sunspot groups is 3.5%. [The determination of the Hale class is done on the (magnetic polarity of the) sunspots, NOT the magnetograms!] 69

  60. Examples from https://www.solarmonitor.org/index.php 70

  61. Figure left: Shin et al. (2016): Development of Daily Maximum Flare-Flux Forecast Models for Strong Solar Flares - http://adsabs.harvard.edu/abs/2016SoPh..291..897S Most of the complex sunspots of the Mount Wilson magnetic classification that are characterized by gamma and/or delta show higher WMFR values (WMFR: weighted mean flare rate). Figure right: Sammis et al. (2000): The Dependence of Large Flare Occurrence on the Magnetic Structure of Sunspots http://adsabs.harvard.edu/abs/2000ApJ...540..583S http://iopscience.iop.org/article/10.1086/309303/pdf In Figure 2, we plot the largest flare from each active region against the largest reported area from that region, for each magnetic class. This shows a roughly linear connection between the logs of SXR flux and active region maximum areas. But the dependence on magnetic class is much stronger. All flares above X4 (4.10-4Wm-2) come from 11 bgd regions of area greater than 1000 MH. Thus, these two conditions constitute a necessary, but not sufficient, condition for an X4 flare. … The general slope of Figure 2, upward and to the right, confirms the well - known fact that large active regions have more large flares than small ones and also tend to be more complex. The probability of a bgd spot group larger than 1000 MH producing an X1 or greater event is only 40%, accounting for about 60% of those events. However, 82% of X1 flares and 100% of the more important X4 events occur in delta spots. By comparison, only 24% of all regions greater than 1000 MH produce X1. The increase in flare size with spot size shows that although the sharp gradient and currents of the delta configuration provide the appropriate situation for flare occurrence, the scale offered by a large spot is important in producing great flares. All large flares (X4 or higher) occur in spot groups of area greater than 1000 MH classified bgd. Predictions that X1 flares will occur for such a class will enjoy a 41% probability of success with no other considerations. Adding some of the considerations mentioned by Zirin & Liggett (1987) and Zirin & Marquette (1991), particularly H-alpha brightness and flux emergence, should improve these predictions considerably. 71

  62. Massi et al.: http://www3.mpifr-bonn.mpg.de/staff/mmassi/c4-Model.pdf  Magnetic shear: the vector magnetic field is oriented parallel to the neutral line than perpendicular to it. Another example of magnetic shear is at MSFC: https://solarscience.msfc.nasa.gov/flares.shtml From https://solarscience.msfc.nasa.gov/flares.shtml Stable sunspots tend to be fairly symmetrical unless there is extensive magnetic shear nearby from emerging magnetic flux or the passing of an area of opposite magnetic polarity. Magnetic shearing can cause large portions of sunspot penumbras to distort or vanish. Lee et al. (2012): Solar Flare Occurrence Rate and Probability in Terms of the Sunspot Classification Supplemented with Sunspot Area and Its Changes http://adsabs.harvard.edu/abs/2012SoPh..281..639L We used sunspot data from 1996 to 2010. We noted that sunspot area and its changes can be a proxy of magnetic flux and its emergence/cancellation, respectively. We classify each sunspot group into the following sub-groups: “Large” and “Small” according to its area and “Decrease”, “Steady”, and “Increase” according to its area changes. Major results from this study can be summarized as follows. i) In the McIntosh sunspot group classification (60 classes in total), the most flare-productive 11 sunspot groups are ‘Dai’, ‘ Eai ’, ‘Fai’, ‘ Dko ’, ‘ Eko ’, ‘ Dki ’, ‘ Dkc ’, ‘ Eki ’, ‘ Ekc ’, ‘ Fki ’, and ‘ Fkc ’. ii) In case of large and compact groups, the flare probabilities noticeably increase with sunspot area. iii) When the sunspot area increases, the flare occurrence rates and probabilities noticeably increase, especially for major flares. Our results show that the sunspot classes having the top five flare occurrence rates are ‘ Fkc ’, ‘ Ekc ’, ‘ Dkc ’, ‘ Fki ’, and ‘ Eki ’, … We note that ‘ Fkc ’ and ‘ Ekc ’ sunspot groups are included in all three studies. This fact may imply that large, asymmetric penumbra sunspot groups should be most flare productive. From the relationship between flare probability and sunspot group area, in the case of large and compact groups, the solar flare probabilities are higher than those of other groups. In the case of “Increase” sub -groups, the flare occurrence rates and probabilities are higher than other sub- groups. This means that when the sunspot area is larger, then the flare probability becomes higher. This is statistical evidence that magnetic flux emergence is an important mechanism for triggering solar flares, because sunspot area can be a good proxy of magnetic flux (Zharkov and Zharkova, 2006). … 72

  63. 7 june 2011 event: STCE: http://www.stce.be/news/353/welcome.html STCE: http://www.stce.be/news/x137x/welcome.html Science at NASA: https://science.nasa.gov/science-news/science-at-nasa/2011/11jul_darkfireworks Superactive regions Chen et al. (2011): Statistical properties of superactive regions during solar cycles 19-23 http://adsabs.harvard.edu/abs/2011A%26A...534A..47C http://www.aanda.org/articles/aa/pdf/2011/10/aa16790-11.pdf Our results indicate that these 45 SARs produced 44% of all the X class X-ray flares during solar cycles 21 – 23, and that all the SARs are likely to produce a very fast CME. The latitudinal distributions of SARs display the Maunder butterfly diagrams and SARs occur preferentially in the maximum period of each solar cycle. Northern hemisphere SARs dominated in solar cycles 19 and 20 and southern hemisphere SARs dominated in solar cycles 21 and 22. In solar cycle 23, however, SARs occurred about equally in each hemisphere. There are two active longitudes in both the northern and southern hemispheres, about 160° – 200° apart and with half-widths of 45°. Criteria to be considered as a SAR: We refer to an AR as a SAR, if three of the four criterion conditions listed in Table 1 are fulfilled. The soft X-ray flare index is the sum of the numerical multipliers of M and X class X-ray flares for the disk transit of the AR. When applying all four criterion conditions, there are 45 SARs selected in solar cycles 21 – 23. Table 1. The four criterion conditions used to parameterize a SAR. Criterion condition Value Maximum sunspot area >1000 MH Flare index >10.0 10.7 cm peak flux >1000 sfu Δ TSI <−0.1% 73

  64. More info on filaments at at http://www.stce.be/news/219/welcome.html More info on polar crown filaments at: - http://solar.physics.montana.edu/wood/99Prom.html - https://science.nasa.gov/science-news/science-at-nasa/2008/17sep_polarcrown 74

  65. Parenti S. (2014): Solar Prominences: Observations http://link.springer.com/article/10.12942/lrsp-2014-1 75

  66. Tlatov et al. (2016): Tilt Angles of Solar Filaments over the Period of 1919 – 2014 http://adsabs.harvard.edu/abs/2016SoPh..291.1115T Mawad et al. (2014): Filaments disappearances in relation to solar flares during the solar cycle 23 http://adsabs.harvard.edu/abs/2015AdSpR..55..696M Hao et al. (2015): Statistical Analysis of Filament Features Based on the Hα Solar Images from 1988 to 2013 by Computer Automated Detection Method http://adsabs.harvard.edu/abs/2015ApJS..221...33H 76

  67. Data taken from: Filippov et al. (2008): Causal relationships between eruptive prominences and coronal mass ejections http://adsabs.harvard.edu/abs/2008AnGeo..26.3025F http://www.ann-geophys.net/26/3025/2008/angeo-26-3025-2008.pdf Zirin (1988): Astrophysics of the Sun 77

  68. From STCE: Two spectacular filament eruptions http://www.stce.be/news/218/welcome.html Some information on material falling back to the solar surface: STCE: A gorgeous filament eruption http://www.stce.be/news/297/welcome.html Later in the event, some of the filament material is seen to fall back down to the surface of the Sun (see image above on the left). These parcels of filament material are trapped on dips in the magnetic field. When the field expands during the eruption, some of the field stretches and the dips disappear. If a magnetic flux tube straightens out, but remains connected to the Sun, the plasma on that tube may slide back to the Sun, and this is what we see in the movie. The filament plasma then causes localized brightenings where it hits the surface, as indicated by an arrow in the image above on the right. 78

  69. The figure was taken from D. A. Hillson: Describing probability: The limitations of natural language http://www.risk-doctor.com/pdf-files/emeamay05.pdf Other (nearly equivalent) terms are used in SWx (SWPC, SIDC): Probability (%) Terminology 0-10 Unlikely 10-25 Small chance 25-50 A chance; Possible 50-75 Likely 75-100 Very likely; Expected As far as we know, there’s no clear terminology consistently applied internal or between the SWx prediction services. 79

  70. The figure was taken from D. A. Hillson: Describing probability: The limitations of natural language http://www.risk-doctor.com/pdf-files/emeamay05.pdf Other (nearly equivalent) terms are used in SWx (SWPC, SIDC,…): Probability (%) Terminology 0-10 Unlikely 10-25 Small chance 25-50 A chance; Possible 50-75 Likely 75-100 Very likely; Expected As far as we know, there’s no clear terminology consistently applied internal or between the SWx prediction services. 80

  71. SWPC ’s « The Weekly » User guide (http://www.swpc.noaa.gov/sites/default/files/images/u2/Usr_guide.pdf ; page 15) Proton Events A proton event starts when the integrated proton flux (5-minute average) rises above a specific threshold for at least three points. The two alert thresholds are (1 pfu = particle / cm-2s-1sr-1): - >10 MeV: ≥10 pfu - >100 MeV: ≥ 1 pfu The time of maximum is the time tag of the 5 minute averaged flux value that has the greatest value. From: http://www.stce.be/news/232/welcome.html The issue here is that proton flares are not considered as separate events if the proton flux (particle energies larger than 10 MeV) at the time of the event is still above the threshold of 10 protons per flux unit (pfu). Example: 6 January 2014 X1 proton flare. From: SESC: https://umbra.nascom.nasa.gov/SEP/ (contains list of all proton events since 1976) Please Note: Proton fluxes are integral 5-minute averages for energies >10 MeV, given in Particle Flux Units (pfu), measured by GOES spacecraft at Geosynchronous orbit: 1 pfu = 1 p / cm -2 sr -1 s -1 . SESC defines the start of a proton event to be the first of 3 consecutive data points with fluxes greater than or equal to 10 pfu. The end of an event is the last time the flux was greater than or equal to 10 pfu. This definition, motivated by SESC customer needs, allows multiple proton flares and/or interplanetary shock proton increases to occur within one SESC proton event. Additional data may be necessary to more completely resolve any individual proton event. SEP definition: from SWPC’s glossary http://www.swpc.noaa.gov/content/space-weather-glossary#Solar%20Energetic%20Particles 81

  72. A very good site on the characteristics between gradual and impulsive events is at SEPEM: http://dev.sepem.oma.be/help/sep_intro.html Papaioannou et al. (2016): Solar flares, coronal mass ejections and solar energetic particle event characteristics http://www.swsc-journal.org/articles/swsc/pdf/2016/01/swsc150076.pdf Reames, 2013: The Two Sources of Solar Energetic Particles http://adsabs.harvard.edu/abs/2013SSRv..175...53R Park et al. , 2015: Study of Solar Energetic Particle Associations with Coronal Extreme-ultraviolet Waves http://iopscience.iop.org/article/10.1088/0004-637X/808/1/3/pdf Impulsive SEP events, having short durations of several hours, are associated with impulsive flares. They are electron rich and have enhanced 3He/4He and Fe/O ratios in contrast to nominal coronal values. Also, they are generally distributed within a narrow propagation cone. Gradual SEP events are associated with gradual X-ray flares and type II and type IV radio emission. They are produced by wide and fast CME-driven shocks (Gopalswamy 2003) and have a broad range of source longitudes (Kahler 1994; Reames 1999). Gradual events show typical coronal abundances and are proton rich in contrast to impulsive events. Many events have characteristics of both gradual and impulsive events due to a combination of both flare- and shock- associated particles (Cane et al. 2006). Surprisingly, some of these appear to have poor magnetic connection to the associated flare sites, suggesting that flare-accelerated particles can be distributed over wide angles in interplanetary space either by efficient cross-field transport in interplanetary space or by ejection of flare particles into an expanding source, for example, a CME shock, near the Sun (Wiedenbeck et al. 2013; Dresing et al. 2014). 82

  73. A very good site on the characteristics between gradual and impulsive events (with more examples) is at SEPEM: http://dev.sepem.oma.be/help/sep_intro.html The cartoons were taken from: Reames (1999): Particle acceleration at the Sun and in the heliosphere http://adsabs.harvard.edu/abs/1999SSRv...90..413R The 15 Sep 2001 event (11 pfu) resulted from an M1 (1N) flare on 11:28UT The 10 Apr 2001 event (355 pfu) resulted from an X2 (3B) flare on 05:26UT. The (halo) CME arrived on 11 April around 15UT, hence the maximum proton flux (near 21UT) coincides with the passage of the ICME. From Reames, 1999: Particle acceleration at the Sun and in the Heliosphere http://link.springer.com/content/pdf/10.1023%2FA%3A1005105831781.pdf Meanwhile, the evidence for two types of events grew. Pallavicini et al. (1977) distinguished impulsive and long-duration (gradual) soft X-ray events; the latter were associated with CMEs (Sheeley et al., 1975). Kahler (1992) has reviewed such differences between flares and CMEs. The connection between these two phenomena and energetic particles in space was made when Cane et al. (1986) found that SEPs associated with the two classes of X-ray events had different proton/electron ratios. The terms ‘gradual’ and ‘impulsive’ have stuck, even though time scales, especially X -ray time scales, poorly resolve those acceleration mechanisms we wish to distinguish. 83

  74. Additional (technical) information on the acceleration process is in Petrosian (2016): Particle acceleration in solar flares and associated CME shocks http://adsabs.harvard.edu/abs/2016ApJ...830...28P In one group, referred to as “prompt” events, the SEPs appear to originate almost simultaneously with the radiations from the flare site located in the lower corona near the tops of reconnecting magnetic loops. The second group, referred to as “delayed” events, shows a complex temporal relation, often with the deduced time of emission of SEPs at the Sun coming after that of HXRs or type III radio. A similar dichotomy seems to be present in the observations of SEP ions. Shorter, weaker events, often referred to as “impulsive,” appear to have higher enrichment of 3He and heavier than CNO ions and softer spectra with an unusual convex spectral shape for enriched ions, while longer-duration and stronger events show near-normal abundances and harder broken power-law spectra. Both the delayed-electron and gradual-normal abundance ion events are more likely to be associated with a fast CME. Here we consider an alternative model where seed particles come from the downstream region. Because of the strong temporal relation, the weak spectral correlation, and consistently harder spectra of the delayed events, we are led to consider a scenario where the SEP spectra result from reacceleration of flare site electrons. This reacceleration is the cause of the harder SEP than HXR- producing electron spectra in these events. An important assumption here is that the CME launch and acceleration of the electrons at the reconnection are almost simultaneous, and/or the upward- escaping accelerated electrons are trapped by turbulence (or other means) in the downstream region of the shock. As described in Section 3.2, we need only a small fraction of these electrons to be trapped and reaccelerated, which renders this assumption reasonable. 84

  75. Animation source: SPACECAST: http://fp7-spacecast.eu/index.php?page=seps Aran A., N. Agueda, C. Jacobs et al. 2011, American Geophysical Union, Fall Meeting 2011, abstract #SH33B-2051A. Solar energetic particles - Modelling of gradual SEP events: an example The movie below shows the modelling of the evolution of the interplanetary shock and the cobpoint (= the point of the shock front that is magnetically connected to the observer), and the fitting to the proton intensity-time profiles measured at 1.0 AU during the gradual proton event on 13 December 2006. The left panel shows: - Three measured proton intensity-time profiles (coloured dots) at L1, the lower energy channel is from the SOHO/ERNE instrument and the other two from the STEREO-A/HET telescope. - The vertical green line marks the shock passage at L1. - The white lines represent the synthetic flux profiles obtained from the fitting using the shock-and-particle model. This model assumes that the shock-accelerated particles are injected onto the interplanetary magnetic field line connecting with the observer at the cobpoint. In this event, high-energy particles are observed shortly after the first cobpoint is established, thus the first injection of shock accelerated particles occurs when the shock is still close to the Sun. The lower energy profiles start increasing later than that because of the presence of a background population of particles, but mainly due to their smaller velocity and due to particle propagation effects along the interplanetary magnetic field. Note that the high energy proton flux peaks a few hours after the onset of the event, while the low-energy intensity peaks at the shock arrival. Thus, indicating that the shock is efficient at accelerating high energy protons when it is close to the Sun, but as it propagates away, it becomes only efficient at accelerating low-energy particles. The >10MeV proton flux as measured by GOES for this event is at ftp://ftp.swpc.noaa.gov/pub/warehouse/2015/2015_plots/proton/20151030_proton.gif 85

  76. Papaioannou et al. (2016): Solar flares, coronal mass ejections and solar energetic particle event characteristics http://www.swsc-journal.org/articles/swsc/pdf/2016/01/swsc150076.pdf Image taken from their Figure 5 (D). Major farside flare from 23 July 2012: http://www.stce.be/news/152/welcome.html Example from 29 October 2015 event: http://www.stce.be/news/325/welcome.html It is believed that the coronal waves associated to strong CMEs widen the access possibilities of the energetic particles to earth coronal Parker spiral magnetic field lines. This would allow some of the backside events still to create a proton event at Earth. But even then, there’s still a problem of the eastern hemisphere events, in particular e.g. at locations E120. Aletrantice mechanism for proton events with farside source: Laitinen et al., 2016: Solar energetic particle access to distant longitudes through turbulent field-line meandering http://www.aanda.org/articles/aa/pdf/2016/07/aa27801-15.pdf We developed a new SEP transport model that takes the non-diffusive propagation of SEPs early in the event history into account for a Parker spiral geometry. We showed that the early onset of SEPs over a wide range of longitudes can be explained by field-line random walk and requires an SEP transport model that properly describes the non-diffusive early phase of SEP cross-field propagation. Our FP+FLRW model is the first model that is capable of reproducing the observed fast access of SEPs to distant longitudes, when the particle and field-line diffusion coefficients are consistently derived from an interplanetary turbulence model. When the FLRW is not included (in the FP model), a much narrower cross-field extent of the SEP event is produced. We conclude that introducing field-line wandering into SEP modelling has the potential of resolving the problem of fast access of SEPs to a wide range of longitudes. (FP: Fokker-Planck; FLRW: Field Lines Random Walk) 86

  77. Papaioannou et al. (2016): Solar flares, coronal mass ejections and solar energetic particle event characteristics http://www.swsc-journal.org/articles/swsc/pdf/2016/01/swsc150076.pdf - Major farside flare from 23 July 2012: http://www.stce.be/news/152/welcome.html - Example from 29 October 2015 event: http://www.stce.be/news/325/welcome.html The likely source of the proton event (NOAA 2434) was 39 degrees behind the west limb as seen from Earth. Mechanism for proton events with farside source: Laitinen et al., 2016: Solar energetic particle access to distant longitudes through turbulent field-line meandering http://www.aanda.org/articles/aa/pdf/2016/07/aa27801-15.pdf We developed a new SEP transport model that takes the non-diffusive propagation of SEPs early in the event history into account for a Parker spiral geometry. We showed that the early onset of SEPs over a wide range of longitudes can be explained by field-line random walk and requires an SEP transport model that properly describes the non-diffusive early phase of SEP cross-field propagation. Our FP+FLRW model is the first model that is capable of reproducing the observed fast access of SEPs to distant longitudes, when the particle and field-line diffusion coefficients are consistently derived from an interplanetary turbulence model. When the FLRW is not included (in the FP model), a much narrower cross-field extent of the SEP event is produced. We conclude that introducing field-line wandering into SEP modelling has the potential of resolving the problem of fast access of SEPs to a wide range of longitudes. (FP: Fokker-Planck; FLRW: Field Lines Random Walk) Aside GOES, there are other satellites equipped with SREM that also measure proton fluxes: http://space-env.esa.int/index.php/SREM_Plots.html Currently (2017), only the Integral satellite has an operational SREM. Finally, STEREO and SOHO also measure proton fluxes: ACE: http://services.swpc.noaa.gov/images/ace-sis-3-day.gif STEREO: https://stereo-ssc.nascom.nasa.gov/beacon/beacon_insitu.shtml (click plots, scroll down to « IMPACT »). 87

  78. Figure taken and slightly modified from Lavraud et al. (2016): A small mission concept to the Sun-Earth Lagrangian L5 point for innovative solar, heliospheric and space weather science http://adsabs.harvard.edu/abs/2016JASTP.146..171L It is likely to have a much faster increase from SEP events from a source on the western hemisphere. When a shock from the associated CME passes the observer (earth), the proton flux will peak (again). For Earth, the solar magnetic footpoint of the Parker spiral associated with a nominal solar wind speed of 400 km/s corresponds to a longitude of about W60. See the graph at http://solar.physics.montana.edu/ypop/Nuggets/2002/021004/parker_spiral.png in the Yohkoh Science nugget at http://solar.physics.montana.edu/ypop/Nuggets/2002/021004/021004.html 88

  79. From the SWPC webpage (http://www.swpc.noaa.gov/noaa-scales-explanation ) NOAA Space Weather Scales The NOAA Space Weather Scales were introduced as a way to communicate to the general public the current and future space weather conditions and their possible effects on people and systems. Many of the SWPC products describe the space environment, but few have described the effects that can be experienced as the result of environmental disturbances. These scales are useful to users of our products and those who are interested in space weather effects. The scales describe the environmental disturbances for three event types: geomagnetic storms, solar radiation storms, and radio blackouts. The scales have numbered levels, analogous to hurricanes, tornadoes, and earthquakes that convey severity. They list possible effects at each level. They also show how often such events happen, and give a measure of the intensity of the physical causes. The « S » stands for Solar radiation Storm. Since observations started in 1976, no S5 event has been recorded. More at http://www.stce.be/news/366/welcome.html 89

  80. From the SWPC webpage (http://www.swpc.noaa.gov/noaa-scales-explanation ): NOAA Space Weather Scales The NOAA Space Weather Scales were introduced as a way to communicate to the general public the current and future space weather conditions and their possible effects on people and systems. Many of the SWPC products describe the space environment, but few have described the effects that can be experienced as the result of environmental disturbances. These scales are useful to users of our products and those who are interested in space weather effects. The scales describe the environmental disturbances for three event types: geomagnetic storms, solar radiation storms, and radio blackouts. The scales have numbered levels, analogous to hurricanes, tornadoes, and earthquakes that convey severity. They list possible effects at each level. They also show how often such events happen, and give a measure of the intensity of the physical causes. The « S » stands for Solar radiation Storm. Since observations started in 1976, no S5 event has been recorded. More on NOAA scales at http://www.stce.be/news/366/welcome.html More on proton intensity at http://www.stce.be/news/233/welcome.html 90

  81. From the SWPC webpage (http://www.swpc.noaa.gov/noaa-scales-explanation ) NOAA Space Weather Scales The NOAA Space Weather Scales were introduced as a way to communicate to the general public the current and future space weather conditions and their possible effects on people and systems. Many of the SWPC products describe the space environment, but few have described the effects that can be experienced as the result of environmental disturbances. These scales are useful to users of our products and those who are interested in space weather effects. The scales describe the environmental disturbances for three event types: geomagnetic storms, solar radiation storms, and radio blackouts. The scales have numbered levels, analogous to hurricanes, tornadoes, and earthquakes that convey severity. They list possible effects at each level. They also show how often such events happen, and give a measure of the intensity of the physical causes. The « S » stands for Solar radiation Storm. Since observations started in 1976, no S5 event has been recorded. More at http://www.stce.be/news/366/welcome.html 91

  82. Bronarska et al. (2017): Characteristics of active regions associated to large solar energetic proton events http://adsabs.harvard.edu/abs/2017AdSpR..59..384B Papaioannou et al. (2016): Solar flares, coronal mass ejections and solar energetic particle event characteristics http://www.swsc-journal.org/articles/swsc/pdf/2016/01/swsc150076.pdf Dierckxsens et al. (2015): Relationship between Solar Energetic Particles and Properties of Flares and CMEs: Statistical Analysis of Solar Cycle 23 Events http://link.springer.com/content/pdf/10.1007%2Fs11207-014-0641-4.pdf Posner et al. (2009): A New Trend in Forecasting Solar Radiation Hazards http://adsabs.harvard.edu/abs/2009SpWea...7.5001P Monitoring of relativistic electrons, which arrive a few to tens of minutes before the actual high- energetic ions. Also in https://arxiv.org/pdf/1210.4475.pdf Example: Bastille day event: https://soho.nascom.nasa.gov/hotshots/2000_07_14/costep_000714.gif 92

  83. Papaioannou et al. (2016): Solar flares, coronal mass ejections and solar energetic particle event characteristics http://www.swsc-journal.org/articles/swsc/pdf/2016/01/swsc150076.pdf Dierckxsens et al., 2015: Relationship between Solar Energetic Particles and Properties of Flares and CMEs: Statistical Analysis of Solar Cycle 23 Events http://link.springer.com/content/pdf/10.1007%2Fs11207-014-0641-4.pdf 93

  84. Papaioannou et al. (2016): Solar flares, coronal mass ejections and solar energetic particle event characteristics http://www.swsc-journal.org/articles/swsc/pdf/2016/01/swsc150076.pdf Dierckxsens et al., 2015: Relationship between Solar Energetic Particles and Properties of Flares and CMEs: Statistical Analysis of Solar Cycle 23 Events http://link.springer.com/content/pdf/10.1007%2Fs11207-014-0641-4.pdf 94

  85. Probabilities and intensities for proton flares for a variety of the above parameters can be found in: Dierckxsens et al., 2015: Relationship between Solar Energetic Particles and Properties of Flares and CMEs: Statistical Analysis of Solar Cycle 23 Events http://link.springer.com/content/pdf/10.1007%2Fs11207-014-0641-4.pdf More info on SEP events at http://dev.sepem.oma.be/help/solpenco2_intro.html 95

  86. Figure to the right taken from https://ase.tufts.edu/cosmos/print_images.asp?id=27 A magnetic reconnection takes place at a current sheet ( dark vertical line ) beneath a prominence and above closed magnetic field lines. The coronal mass ejection, abbreviated CME, traps hot plasma below it ( hatched region ). The solid curve at the top is the bow shock driven by the CME. The closed field region above the prominence ( center ) is supposed to become a flux rope in the interplanetary medium. [Adapted from Petrus C. Martens and N. Paul Kuin (1989).] In this model of a three-part coronal mass ejection, portrayed by Terry Forbes (2000), swept-up, compressed mass and a bow shock have been added to the eruptive-flare portrayal of Tadashi Hirayama (1974). The combined representation includes compressed material at the leading edge of a low-density, magnetic bubble or cavity, and dense prominence gas. The prominence and its surrounding cavity rise through the lower corona, followed by sequential magnetic reconnection and the formation of flare ribbons at the footpoints of a loop arcade. [Adapted from Hugh S. Hudson, Jean- Louis Bougeret and Joan Burkepile (2006).] Figure to the left taken from Forbes (2000): A review on the genesis of coronal mass ejection http://adsabs.harvard.edu/abs/2000JGR...10523153F http://onlinelibrary.wiley.com/doi/10.1029/2000JA000005/epdf When CMEs were first clearly identified by Skylab in 1973, many researchers assumed that they were caused by the outward expansion of hot plasma produced by a large flare. We now know that this is not the case, for several reasons. First, less than 20% of all CMEs are associated with large flares [Gosling, 1993]. Second, CMEs that are associated with flares often appear to start before the onset of the flare [Wagner et al., 1981; Simnett and Harrison, 1985]. Finally, the thermal pressure produced by a flare is too small to blow open the strong magnetic field of the corona. 96

  87. Expanding flux rope consisting of magnetic field lines and filament at bottom (left sketch) evolves into a CME with leading edge (pushed up by top of magnetic field lines from flux rope), cavity and core filament (right sketch). The shock is not visible, but measurable by e.g. Type II radio bursts. Source file: Webb et al. (2012): Coronal Mass Ejections: Observations http://link.springer.com/article/10.12942/lrsp-2012-3 However, with the larger dynamic range of LASCO rims of material detected ahead of fast LASCO CMEs are now considered evidence of shock waves, and emission can be detected ahead of slower speed CMEs as low-level brightness enhancements due to the expanding streamer (see Section 3.6). The interrelationship between the various features which one can associate with CMEs is shown in Figure 3. It should be kept in mind that these features are not necessarily present in all CMEs. Not all CMEs contain a prominence, nor do all CMEs have detectable chromospheric ribbons and shock waves * Webb et al. (2012): Coronal Mass Ejections: Observations http://link.springer.com/article/10.12942/lrsp-2012-3 The original definition of a CME as a new, discrete brightening in the field of view over a time-scale of tens of minutes which is always observed to move outward (e.g., Webb and Hundhausen, 1987) is still generally accepted. However, some workers tend to regard any eruption from the Sun observed in the corona, no matter how faint or narrow, as a CME while others regard an eruption as a CME only if it has a certain size or structure. Although a “typical” CME is now thought to involve the eruption of a magnetic flux rope, the structure and magnitude of any CME magnetic field near the Sun can only be inferred, since we cannot directly measure coronal magnetic fields. 97

  88. Expanding flux rope consisting of magnetic field lines and filament at bottom (left sketch) evolves into a CME with leading edge (pushed up by top of magnetic field lines from flux rope), cavity and core filament (right sketch). The shock is not visible, but measurable by e.g. Type II radio bursts. Source file: Webb et al. (2012): Coronal Mass Ejections: Observations http://link.springer.com/article/10.12942/lrsp-2012-3 However, with the larger dynamic range of LASCO rims of material detected ahead of fast LASCO CMEs are now considered evidence of shock waves, and emission can be detected ahead of slower speed CMEs as low-level brightness enhancements due to the expanding streamer (see Section 3.6). The interrelationship between the various features which one can associate with CMEs is shown in Figure 3. It should be kept in mind that these features are not necessarily present in all CMEs. Not all CMEs contain a prominence, nor do all CMEs have detectable chromospheric ribbons and shock waves * Webb et al. (2012): Coronal Mass Ejections: Observations http://link.springer.com/article/10.12942/lrsp-2012-3 The original definition of a CME as a new, discrete brightening in the field of view over a time-scale of tens of minutes which is always observed to move outward (e.g., Webb and Hundhausen, 1987) is still generally accepted. However, some workers tend to regard any eruption from the Sun observed in the corona, no matter how faint or narrow, as a CME while others regard an eruption as a CME only if it has a certain size or structure. Although a “typical” CME is now thought to involve the eruption of a magnetic flux rope, the structure and magnitude of any CME magnetic field near the Sun can only be inferred, since we cannot directly measure coronal magnetic fields. 98

  89. Expanding flux rope consisting of magnetic field lines and filament at bottom evolves into a CME with leading edge (pushed up by top of magnetic field lines from flux rope), cavity and core filament. Left picture: SOHO Gallery: https://sohowww.nascom.nasa.gov/gallery/images/las02.html Right picture: STCE: http://www.stce.be/news/342/welcome.html Coronagraph: Wiki:a telescopic attachment designed to block out the direct light from the Sun so that nearby objects – which otherwise would be hidden in the star's bright glare – can be resolved. In short: it is an instrument to create a permanent total solar eclipse. Coronagraph Lasco: https://lasco-www.nrl.navy.mil/index.php?p=content/handbook/hndbk_5 CMEs are mostly observed in white light by coronagraphs from space (SOHO, STEREO). In order to make the faint CMEs better visible, difference images are used (one image subtracted from the other). Ground-based observatories can observe CMEs very close to the Sun: MLSO (K-Cor): http://download.hao.ucar.edu/d5/www/fullres/latest/latest.kcor.gif Ground-based observatories can also observe CMEs by using interplanetary scintillation (IPS). * Dorrian et al. (2008): Simultaneous interplanetary scintillation and Heliospheric Imager observations of a coronal mass ejection https://core.ac.uk/download/pdf/16283575.pdf Interplanetary scintillation (IPS) was first described by Hewish et al. [1964]. When the raypath from a compact radio source passes through the solar wind it encounters regions of varying plasma density, inducing phase variations. As the wave continues to the receiver these phase variations are converted into amplitude variations by interference [e.g., Coles, 1978]. 99

  90. Source file: Webb et al. (2012): Coronal Mass Ejections: Observations http://link.springer.com/article/10.12942/lrsp-2012-3 * Webb et al. (2012): Coronal Mass Ejections: Observations http://link.springer.com/article/10.12942/lrsp-2012-3 5.2 Interplanetary scintillation (IPS) observations The IPS technique relies on measurements of the fluctuating intensity level of a large number of point- like distant meter-wavelength radio sources. They are observed with one or more ground arrays operating in the MHz – GHz range. IPS arrays detect changes to density in the (local) interplanetary medium moving across the line of sight to the source. Disturbances are detected by either an enhancement of the scintillation level and/or an increase in velocity. When built up over a large number of radio sources a map of the density enhancement across the sky can be produced. The technique suffers from relatively poor temporal (24-hour) resolution and has a spatial resolution limited to the field of view of the radio telescope. For example, high-latitude arrays such as the long- deactivated 3.5 ha array near Cambridge in the UK could not observe sources in the mid-high latitude southern hemisphere. Scattering efficiency also poses a limitation on IPS measurements as increasing the frequency at which to measure the sources allows an observer to detect disturbances closer to the Sun. Higher frequencies means fewer sources, however, so the spatial resolution is effectively decreased. Finally, ionospheric noise limits viewing near the Sun and near the horizon, and a model- dependence for interpreting the signal as density or mass. Workers have, however, been working with these difficulties for 50 years and a number of techniques have evolved to extract reliable CME measurements using IPS. Recent papers involving such measurements include Jones et al. (2007), Bisi et al. (2008), Jackson et al. (2010b), Tappin and Howard (2010), and Manoharan (2010). 100

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